A. Cotera, S. Markoff, T. R. Geballe, and H. Falcke (Eds.)
Proceedings of the
Galactic Center Workshop 2002 The Central 300 parsecs of the Milky Way
Astronomische Nachrichten, Supplementary Issue 112003
WILEYVCH WILEY-VCH Verlag CmbH & Co. KCaA
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A. Cotera, S. Markoff, T. R. Geballe, and H. Falcke (Eds.)
Galactic Center Workshop 2002 The Central 300 parsecs of the Milky Way Astronomische Nachrichten, Supplementary Issue 112003
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A. Cotera, S. Markoff, T. R. Geballe, and H. Falcke (Eds.)
Proceedings of the
Galactic Center Workshop 2002 The Central 300 parsecs of the Milky Way
Astronomische Nachrichten, Supplementary Issue 112003
WILEYVCH WILEY-VCH Verlag CmbH & Co. KCaA
Editors DI:Angela Cotera SETI Institute, Arizona State University
[email protected] This book was carefully produced. Nevertheless, editors, authors and publisher do not warrant the information contained therein to be free of errors. Readers are advised to keep in mind that statements, data, illustrations, procedural details or other items may inadvertently be inaccurate.
DI:Sera Markoff Max Planck Institute, Center for Space Research smarkoff @ mpifr-bonn.mpg.de Pro$ DI:Heino Falcke Max Planck Institute, Center for Space Research
[email protected] Thoinas R. Geballe Gemini Observatory
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Preface
The Galactic Center (GC) is one of the most scientifically intriguing regions available for astrophysical research. First and foremost, it is by far the closest galactic nucleus available for study, observable at spatial resolutions unapproachable in other galaxies. At the distance of the GC (-8 kpc), spatial resolutions of %0.1”, now achievable with the latest generation of radio and near-infrared telescopes, corresponds to s 800 AU. The highest achievable resolution for the closest external galactic nucleus, M31, is a factor of 100 times lower. Thus the relative proximity of our Galactic Center, combined with the high angular resolutions now available, provides the opportunity to differentiate and study a wealth of objects; many of which appear to be highly unusual, possibly unique, but could in fact be merely the standard occupants of a normal galactic nucleus. At the very center of the Galaxy lurks the (currently) weakly luminous, supermassive black hole candidate Sgr A*. Surrounding this object is the densest stellar cluster in the Galaxy, much of it immersed in ionized interstellar gas. At much greater scales of up to 100 pc, clusters of bright stars and synchrotron-brightfilaments are prominent. All of these phenomena are surrounded by the galactic stellar bulge and dense giant molecular clouds extending out to kiloparsec scales. Proximity to the GaIactic Center does not, however, make observations of this region trivial. Due to large amounts of dust and gas along the line of sight to the GC - through the Galactic disk - the central region is inaccessible with optical astronomy, even using the most the sensitive and sophisticated techniques. Observational exploration of the GC had to await, and must rely on, the development of radio, sub-millimeter, infrared, and most recently, X-ray and y-ray astronomy. The first subarcsecond observations of the Galactic Center were made in the 1970’s with radio interferometers. In the 1980’s the VLA led the way in high angular resolution studies, while infrared observations began to reach the sub-arcsecond regime. Technological developments in infrared observing with both ground- and space-based telescopes, as well as an improved understanding of atmospheric seeing, have allowed significantly higher resolution observations of the GC in the near-infrared (NIR, 1-3 pm) and mid-infrared (MIR, 3-20 pm). More recently, remarkable improvements in resolution at far-infared (FIR, 20-200 pm), sub-millimeter, millimeter, and within the last few years, X-ray, wavelengths, have resulted in remarkable progress in addressing many long standing questions about the Galactic Center. For example, studies of stellar orbital motions have confirmed that the GC harbors what is by far the most secure candidate for a supermassive (-3 x lo6 M,) black hole: Sgr A*. In very recent years, the detection of weak X-ray emission undergoing daily flares has prompted numerous revisions to theoretical models of the accretion processes onto this object. The Galactic Center also contains a collection of some of the most luminous stars in the Galaxy, which are responsible for both compact and diffuse X-ray emission and the ioniza-
VI
Preface
Preface
VII
tion of large quantities of interstellar gas. The obvious youth of these stars raises questions about how close to the center they formed, and how they formed in an environment that currently appears to be quite hostile to starbirth. These luminous stars ionize the surfaces of regions of molecular gas, some of which appear to be interacting with both the hot gas and strong ambient magnetic fields. Thus, the Galactic Center is complex laboratory containing a multitude astrophysical phenomena. Currently the links between some of these are obvious, and between others are less so. The detailed investigation of the GC is fraught with exciting questions, whose answers have ramifications for the structure and evolution of our galaxy, the phenomena seen in more distant galactic nuclei, and even fundamental physics. The study of the GC brings together many disciplines of astronomy which are tycpically distinct, and many astronomers whose research does not ordinarily overlap. The Galactic Center Workshop 2002: The Central 300 Parsecs, follows a series of conferences which started in the 1970’s, dedicated specifically to the center of our Galaxy. From one and two day symposia at Caltech and the University of California at Berkeley (1982, 1984, and 1986), these meetings have continued to expand, with the fiist week-long conference held in 1988 at the University of California at Los Angeles. The meetings have also become more international, with the 1992 meeting in Ringsburg, Germany, the 1995 meeting in La Serena, Chile, as well as a symposium during the 1997 IAU General Assembly meeting in Kyoto, Japan. In 1998, The Galactic Center Workshop 1998: The Central Parsecs, held in Tucson, Arizona, was the first full meeting dedicated solely to the very heart of our Galaxy. The Galactic Center Workshop 2002: The Central 300 parsecs, held at the Ohana Keauhou Beach Resort in Kona, Hawai’i on 3-8 November, consisted of eight topical sessions. The format of each of these was 6-8 talks followed by 30-45 minutes of moderated discussion, which was recorded. These Proceedings are arranged in the same order as the talks, with the posters from each session following. The discussions were informative and often lively; a CD of the recorded discussions is also being made available as part of these Proceedings. This current workshop was extremely fortunate to have been sponsored by so many prestigious institutions, most of them Hawai’i observatories. The Gemini Observatory was the primary host of the meeting. Significant financial support was provided by Gemini, the Institute for Astronomy, University of Hawaii; the Subaru Telescope; the Joint Astronomy Centre; the W. M. Keck Observatory; the Sub-Millimeter Array; and the National Science Foundation. The Galactic Center Newsletter played a key role in disseminating information about the meeting to astronomers worldwide. As conference co-chairs, we are grateful to the more than 100 attendees for their enthusiastic response to this conference. We would like to especially thank the tireless work of the local organizing committee, without whom “GC02” would not have been such a success. Angela Cotera Sera Markoff Thomas Geballe Heino Falcke
VIII
Scientitic Organizing Committee Angela Cotera (co-chair), Anzona State University Andreas Eckart, University of Cologne Heino Falcke (co-chair), Max Planck Institute, Center for Space Research Tom Geballe (co-chair), Gemini Observatory Rolf Kudritzki, Institute for Astronomy, Honolulu Sera Markoff (co-chair), Max Planck Institute, Center for Space Research Mark Moms, University of California at Los Angeles Ramesh Narayan, Harvard University Torno Oka, University of Tokyo Mark Reid, Harvard University Francois Rigaut, Gemini Observatory Kris Sellgren, Ohio State University Farhad Yusef-Zadeh, Northwestern University Jun-Hui Zhao, Harvard-Smithsonian Center for Astrophysics
Local Organizing Committee Andy Adamson, Joint Astronomy Centre Tracy Beck, Gemini Observatory Tom Geballe (chair), Gemini Observatory John Hamilton, Gemini Observatory Janice Harvey, Gemini Observatory Kalena Quinones, Gemini Observatory Peter Michaud, Gemini Observatory Francois Rigaut, Gemini Observatory Antony Schinckel, Sub-Millimeter Array
Preface
Contents
Chapter 1: Large Scale Structures: Surveys and Interactions High-resolution Hi Absorption Observations towards the Central 200 pc of the Galaxy Cornelia C. Lang, Claudia Cyganowski, W. M. Goss, Jun-Hui Zhao
.......................................................
High Resolution, High Sensitivity Imaging of the Galactic Center at 330 MHz Michael E. Nord, Crystal L. Brogan, Scott D. Hyman, T. Joseph W. Lazio, Namir E. Kassim, T.N. LaRosa, K. Anantharamaiah, Neboja Duric
..
Spatially Resolved Very Large Array 74 MHz Observations Toward the Galactic 17 Center C. L. Brogan, M. Nord, N. Kassim, J. Lazio, K. Anantharamaiah
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Chandra view of the central 300 pc of our Galaxy
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25
Q. Daniel Wang
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Two Thousand X-ray Stars in the Central 20 pc of the Galaxy 33 M. P. Muno, F. K. Baganoff, M. W. Bautz, W. N. Brandt, P. S. Broos, E. D. Feigelson, G. P. Garmire, M. R. Morris, G. R. Ricker, L. K. Townsley Magnetic field in the Galactic Centre: Rotation Measure observations of extra41 galactic sources Subhashis Roy, A. Pramesh Rao, Ravi Subrahmanyan
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47
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53
Study of the Nuclear Bulge region of the Galaxy K. S. Baliyan, S. Ganesh, U. C. Joshi, I. S. Glass, T. Nagata
A morphological Study of the Galactic Inner Bulge Kiran S. Baliyan, Shashikiran Ganesh, Umesh C. Joshi, Ian S. Glass, Mark R. Morris, Alain Omont, Mathias Schultheis, Guy Simon Warm molecular gas, dust and ionized gas in the 500 central pc of the Galaxy N. J. Rodriguez-FernBndez, J. Martin-Pintado, A. Fuente, T. L. Wilson
..
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Prospects for LOFAR Observations of the Galactic Center N. E. Kassim, T. J. W. Lazio, M. Nord, S. D. Hyman, C. L. Brogan, T. N. LaRosa, N. Duric
59
65
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Contents
43 GHz SiO masers in late-type stars with 86 GHz SiO masers and astrometry 73 with VERA in the Galactic center LorAnt 0. Sjouwerman, Maria Messineo, Harm J. Habing
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A Search for Radio Transients at 0.33 GHz in the GC Scott D. Hyman, T. Joseph W. Lazio, Namir E. Kassim, Michael E. Nord, Jennifer L. Neureuther
79
Chapter 2: Molecular Clouds and Magnetic Fields
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A Molecular Face-on View of the Galactic Center Region Tsuyoshi Sawada, Tetsuo Hasegawa, Toshihiro Handa, R. J. Cohen
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The Inner 200pc: Hot Dense Gas Christopher L. Martin, Wilfred M. Walsh, Kecheng Xiao, Adair P. Lane, Christopher K. Walker, Antony A. Stark Gravitational Stability of Molecular Clouds in the Galactic Center Tomoharu Oka, Tetsuo Hasegawa
85
93
........... 101
Spectroscopy of Hydrocarbon Grains toward the Galactic Center and Quintuplet Cluster 109 J. E. Chiar, A. J. Adamson, D. C. B. Whittet, Y .J. Pendleton
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X-rays from the HII Regions and Molecular Clouds near the Galactic Center Katsuji Koyama, Hiroshi Murakami, Shinichiro Takagi
. . 117
Reflected X-ray Emissions on Giant Molecular Clouds - Evidence of the Past Activities of Sgr A* 125 Hiroshi Murakami, Atsushi Senda, Yoshitomo Maeda, Katsuji Koyama
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Observation of Toroidal Magnetic Fields on 100 pc Scales in the Galactic Center 133 G. Novak, D. T. Chuss, J. L. Dotson, G. S. Griffin, R. F. Loewenstein, M. G . Newcomb, D. Pernic, J. B. Peterson, T. Renbarger Extended photoionization and photodissociation in Sgr B2 J. R. Goicoechea, N. J. Rodriguez-Fern6ndez, J. Cernicharo Propagation of charged particles from the Galactic Center W. Bednarek, M. Giller, M. Zielifiska
................. 139
................. 145
Discovery of New SNR Candidates in the Galactic Center Region with ASCA and Chandra 151 Atsushi Senda, Hiroshi Murakami, Katsuji Koyama
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Contens
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Molecular Lineobservations of the Tornado Nebula and its Eye J. Lazendic, M. Burton, F. Yusef-Zadeh, M. Wardle, A. Green, J. Whiteoak
The Search for Water and Other Molecules in the Galactic Centre with the Odin Satellite 161 Aa. Sandqvist, P. Bergman, A. Hjalmarson, E. Falgarone, T. Liljestrom, M. Lindqvist, A. Winnberg, the Odin Team
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Chapter 3: Sagittarius A and its Environs
... 167
Deep X-Ray Imaging of the Central 20 Parsecs of the Galaxy with Chandra Mark Morris, Fred Baganoff, Michael Muno, Christian Howard, Yoshitomo Maeda, Eric Feigelson, Marshall Bautz, Niel Brandt, George Chartas, Gordon Garmire, Lisa Townsley
........... 173
Mapping Magnetic Fields in the Cold Dust at the Galactic Center David T. Chuss, Giles Novak, Jacqueline A. Davidson, Jessie L. Dotson, C. Darren Dowell, Roger H. Hildebrand, John E. Vaillancourt
The Galactic Center Nonthermal Filaments: Recent Observations and Theory T. N. LaRosa, Michael E. Nord, T. Joseph W. Lazio, Steven N. Shore, Namir E. Kassim
. . 181
Interaction between the Northeastern Boundary of Sgr A East and Giant Molecu189 lar Clouds: Excitation Mechanisms of the H2 Emission Sungho Lee, Soojong Pak, Christopher J. Davis, Robeson M. Hermstein, T. R. Geballe, Paul T. P. Ho, J. Craig Wheeler
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Sgr A East and its surroundings - a view with XMM-Newton Masaaki Sakano, Robert S. Warwick, Anne Decourchelle
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211
Chandra ACIS Imaging Spectroscopy of Sgr A East Y. Maeda, K. Itoh, F. K. Baganoff, M. W. Bautz, W. N. Brandt, D. N. Burrows, J. P. Doty, E. D. Feigelson, G. P. Garmire, M. Morris, M. P. Muno, S. Park, S. H. Ravdo, G. R. Ricker, L. K. Townsley A Census of Dust Absorption at the Galactic Centre Andy Adamson, Rachel Mason, Emily Macdonald, Gillian Wright, Jean Chiar, Yvonne Pendleton, Tom Kerr, Janet Bowey, Doug Whittet, Mark Rawlings
Thermal SiO observations of a shell attached to the nonthermal filaments in 217 SgrA Toshihiro Handa, Masaaki Sakano, Masato Tsuboi
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Absorption and Emission in the Four Ground-State OH Lines Observed at 18 cm 223 with the VLA Towards the Galactic Centre R. Karlsson, Aa. Sandqvist, L. 0. Sjouwerman, J. B. Whiteoak
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Contents
Constraints on distances to Galactic Centre non-thermal filaments from HI 229 absorption Subhashis Roy
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........... 235
Discovery of a non-thermal X-ray filament in the Galactic Centre Masaaki Sakano, Robert S. Warwick, Anne Decourchelle High-negative velocities in the inner 25 pc of the Galactic center Lorint 0. Sjouwerman
............ 241
Chapter 4: Stars and Star Formation Really Cool Stars and the Star Formation History at the Galactic Center Robert D. Blum, Solange V. Rm’rez, Kristen Sellgren, Knut Olsen Massive Stars and The Creation of our Galactic Center Donald F. Figer
...... 247
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The Galactic Center Source IRS 13E: a Star Cluster Jean-Pierre Maillard, Thibaut Paumard, Susan Stolovy, Franqois Rigaut X-ray Emission from Stellar Clusters Near the Galactic Center Casey Law, Farhad Yusef-Zadeh
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Simulated X-ray emission from the Arches cluster Pablo F. Velrizquez, Alejandro C. Raga, Jorge Cant6, Elena Masciadri, Luis F. Rodriguez
SiO Maser Sources within 30 pc of the Galactic Center Shuji Deguchi, Hiroshi Imai
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86 GHz SiO masing late-type stars in the Inner Galaxy M. Messineo, H. J. Habing, L. 0. Sjouwennan, K. M. Menten, A. Omont CNO Abundances in the Quintuplet Cluster M Supergiant 5-7 S. V. Ramirez, K. Sellgren, R. Blum, D. M. Terndrup
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303
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315
New results on the Galactic Center Helium stars Thibaut Paumard, Jean-Pierre Maillard, Susan Stolovy
Ten Thousand Stars Toward the Galactic Center Franqois Rigaut, Robert Blum, Tim Davidge, Angela Cotera
Stellar Orbits at the Center of the Milky Way N. Mouawad, A. Eckart, S. Pfalzner, J. Moultaka, C. Straubmeier, R. Spurzem, R. Schodel, T. Ott
XI11
Contens
Dynamical Friciton near the Galactic Center Sungsoo S. Kim, Donald F. Figer, Mark Moms
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321
Near-infrared adaptive optics observations of the Galactic Center with NAOS/CONICA (ESO) and GriF (CFHT) 327 Y. Clknet, D. Rouan, F. Lacombe, E. Gendron, D. Gratadour
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Radio Pulsars in the Galactic Center T. Joseph W. Lazio, James M. Cordes, Cornelia C. Lang, Eric V. Gotthelf, Q. Daniel Wang Review of low-mass X-ray binaries near the Galactic center A. Lutovinov, S. Grebenev, S. Molkov, R. Sunyaev
333
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Neutrons, neutrinos, and gamma-rays from the Galactic Center W. Bednarek
............. 343
Chapter 5: Sgr A* I: New Observational Results
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349
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355
Linear and Circular Polarization from G. C. Bower Intrinsic Radio Variability of Sgr A* Jun-Hui Zhao
Flares of Sagittarius A* at Short Millimeter Wavelengths Atsushi Miyazaki, Takahiro Tsutsumi, Masato Tsuboi
.................. 363
Limits on the Short Term Variability of Sagittarius A* in the Near-Infrared S. D. Hornstein, A.M. Ghez, A. Tanner, M. Morris, E. E. Becklin
.... 371
A New X-Ray Flare from the Galactic Nucleus Detected with XMM-Newton A. Goldwum, E. Brion, P. Goldoni, P. Fernando, F. Daigne, A. Decourchelle, R. S. Warwick, P. Predehl
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Searching for Structural Variability in Sgr A* Zhi-Qiang Shen, M. C. Liang, K. Y. Lo, M. Miyoshi
Observations of the Galactic Centre at 610 MHz with the GMRT Subhashis Roy, A. F’ramesh Rao
... 377 383
........... 391
Closure Amplitude Analysis of 15, 22 and 43GHz VLBA Observations of Sagit397 tarius A*: Size is Consistent with the Scattering Law G. C. Bower
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Contents
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VLBA observation of a radio intraday flare of Sgr A* Makoto Miyoshi, Hiroshi Imai, Junichi Nakashima, Shuji Deguchi, Zhi-Qiang Shen
403
A Chandra View of Diffuse X-Ray Emission in the Central 20 Parsecs of the Galaxy 407 Sangwook Park, Frederick K. Baganoff, Mark W. Bautz, Gordon P. Garmire, Yoshitomo Maeda, Mark Morris, Michael P. Muno
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Investigating the submillimetre variability of Sagittarius A* with SCUBA Douglas Pierce-Price, Tim Jenness, John Richer, Jane Greaves
...... 413
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Near-Infrared Flux Limits for Sgr A* Based on NICMOS Data Susan Stolovy, Fulvio Melia, Donald McCarthy, Farhad Yusef-Zadeh
The wavelength dependence of Sgr A* size and the unified model of compact radio 425 sources Fedor Prigara
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Search for Circular Polarization toward Sagittarius A* at 100 GHz M. Tsuboi, H. Miyahara, R. Nomura, T. Kasuga, A. Miyazalu
.......... 431
Chapter 6: Sgr A* 11: Theoretical Models Radiatively Inefficient Accretion Flow Models of Sgr A* Eliot Quataert
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445
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453
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459
Jet Models for Flaring in Sgr A* Sera Markoff, Heino Falcke A Jet-ADAF Model for Sgr A* F. Yuan, S. Markoff, H. Falcke
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A model for polarised radio emission from Sgr A* T. Beckert
On the Chandra Detection of Diffuse X-Ray Emission from Sgr A* M. E. Pessah, F. Melia A Relativistic Disk in Sagittarius A* Siming Liu, Fulvio Melia
.......... 467
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The frozen (inactive) disk in Sgr A*: freezing the accretion of the hot gas too? Sergei Nayakshin Gamma-ray emission from an ADAF around a Kerr black hole Kazutaka Oka, Tadahiro Manmoto
475
. . 483
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Contens
Chapter 7: The Central Parsec
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497
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505
The Discovery of Sgr A* W. M. Goss, Robert L. Brown, K. Y. Lo
The Position, Motion, and Mass of Sgr A* Mark J. Reid, Karl M. Menten, Reinhard Genzel, Thomas Ott, Rainer Schodel, Andreas Brunthaler Tidal processes very near the black hole in the Galactic Center Tal Alexander
............. 513
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New MIR Excess Sources north of the IRS 13 Complex 521 A. Eckart, J. Moultaka, T. Viehmann, C. Straubmeier, N. Mouawad, R. Genzel, T. Ott, R. Schodel Full Three Dimensional Orbits For Multiple Stars on Close Approaches to the 527 Central Supermassive Black Hole A. M. Ghez, E. Becklin, G. DuchCne, S. Hornstein, M. Moms, S. Salim, A. Tanner
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The Galactic Center stellar cluster: The central arcsecond R. Schodel, R. Genzel, T. Ott, A. Eckart
................. 535
Stellar Dynamics in the Galactic Center: 1000 Stars in 100 Nights Thomas Ott, Reinhard Genzel, Andreas Eckart, Rainer Schodel A Bow Shock of Heated Dust Surrounding IRS 8 F. Rigaut, T. R. Geballe, J.-R. Roy, B. T. Draine
........... 543
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551
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Monitoring Sagittarius A* in the MIR with the VLT 557 A. Eckart, J. Moultaka, T. Viehmann, C. Straubmeier, N. Mouawad, R. Genzel, T. Ott, R. Schodel, F. K. Baganoff, M. R. Morris The magnetic field in the central parsec A. C. H. Glasse, D. K. Aitken, P. F, Roche
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563
Mid-Infrared Imaging and Spectroscopic Observations of the Galactic Center with 567 SubardCOMICS Y. Okada, T. Onaka, T. Miyata, H. Kataza, Y. K. Okamoto, S. Sako, M. Honda, T. Yamashita, T. Fujiyoshi
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Physical Conditions in the Central Parsec Modeled from Mid-Infrared Imaging 573 Photometry Dan Gezari, Eli Dwek, Frank Varosi
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LINCINIRVANA - The LBT Near-Infrared Interferometric Camera C. Straubmeier, A. Eckart, T. Bertram, T. Herbst
......... 577
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Contents
Chapter 8: Morphology and Dynamics of the Central 10 Parsecs Hot Molecular Gas in the Central 10 Parsecs of the Galaxy R. M. Herrnstein, P. T. P. Ho
................ 583
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The ISM and Stellar Distributions Near Sgr A* Nick Scoville, Susan R. Stolovy, Micol Christopher
Resolving The Northern Arm Sources at the Galactic Center Angelle M. Tanner, A. M. Ghez, M. Morris, E. E. Becklin
............... 597
Structural analysis of the Minispiral from high-resolution Br/ data Thibaut Paumard, Jean-Pierre Maillard, Mark Moms Gas physics and dynamics in the central 50 pc of the Galaxy B. Vollmer, W. J. Duschl, R. Zylka
591
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The First Measurement of Radial Acceleration of Ionized Gas Near Sagittarius A* 621 Doug Roberts, Farhad Yusef-Zadeh
............. 629
Simple hydrodynamical Simulations of the Circumnuclear Disk Robert F. Coker, Michol H. Christopher, Susan R. Stolovy, Nick Z. Scoville
Astron. Nachr./AN 324. No. S I . 1 - 7 (2003) / DO1 I0.1002/asna.200385065
High-resolution HI Absorption Observations towards the Central 200 pc of the Galaxy Cornelia C. Lang*’, Claudia Cyganowski**2,W. M. G O S Sand ~ , Jun-Hui Zhao’ I Department of Physics & Astronomy. University of Iowa, Iowa City, IA 52242 * National Radio Astronomy Observatory, Socorro, NM 87801 Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Camhndge, MA 01238
’
Key words interstellar medium, Galactic Centet
-
Abstract. We report the first results of an HI absorption survey of the central 200 pc of the Galaxy. These Very Large Array (VLA) data have a resolution of 15” (0.6 pc at the Galactic Center (GC) distance) and a velocity resolution of -2.5 !an s-’ . This study provides HI data with high spatial resolution, comparable with the many high resolution observations which have been made of GC sources over the past ten years. Information on the velocities, relative di\tunces of Sources and H I column densities will be studied. Such data serves to clarify the nature of physical associations between the unique radio continuum features in this region and the atomic and molecular components. In particular, we are interested in the arrangement of sources in the active Radio Arc region: the atomic, ionized and molecular gas and their relation to the massive stellar clusters and magnetic filaments.
1 Introduction The bright and unusual radio continuum sources in the central few hundred parsecs of the Galaxy provide an opportunity to measure the 21 cm line of atomic hydrogen in absorption. Previous HI absorption studies have been crucial to our understanding of Galactic structure and rotation and the nature of atomic gas in the inner parts of the Galaxy. There are several well-known HI components in the direction of the GC (Cohen & Davies 1979): ( 1 ) the expanding 3 kpc arm appears at v=-53 km s-l and is thought to be -5.5 kpc from the Sun, (2) HI components near v=+l35 km s-l are thought to be beyond the G C by distances of a few hundred parsecs to 2 kpc, and (3) the “nuclear disk” and “molecular ring” components appearing at velocities of 160 to -200 km s-’ and 135 km spl are located within a few hundred parsecs of the GC. In the inner Galaxy, atomic gas is most often associated with regions of molecular gas where it serves to shield the molecular gas against photodissociation (Dickey & Lockman 1990). Therefore, HI absorption features not described above may possibly be identified with known G C molecular emission features using correlations in velocity structure. The recent Oka et al. (1998) CO survey made with the Nobeyama 45-m telescope provides the best spatial resolution, velocity, and spatial coverage of any survey of molecular gas within the central Galaxy. In addition, the multitude of “forbidden” (e.g. sign opposite to galactic rotation) velocity components in the GC region are thought to represent the response of the molecular gas in the GC to the Galaxy’s strong stellar bar (Binney et a]. 1991 ; Bally et a]. 1988). The CO survey data of Oka et al. (1998) illustrated that the molecular gas traced by C O emission in the central 200 pc is organized into filamentary and shell-like features. This morphology and kinetic structure indicates that violent kinetic activity (such as supernova explosions and stellar winds from Wolf-Rayet type ~
~
* Corresponding author: e-rnai1:
[email protected] * * C.Cyganowski was a REU summer research student at NRAO @ 2003 WILEY~VCHVerlag GinbH & Co. KGaA. Weinhem
C. Lang et al.: HI towards the Galactic Center
2
stars) plays an important role in shaping the ISM. In addition to the Radio Arc region (where the Quintuplet and Arches clusters are located), the GC region is filled with sites where compact thermal radio and midinfrared sources have been observed (e.g. Sgr B, Sgr C; as well at many positions along the Galactic plane; LaRosa et al. 2000; Egan et al. 1998) and it is likely that massive stars are either forming or have formed in these regions. In addition, the spectrum of diffuse X-ray emission in this region suggests that the ISM is being strongly influenced by massive star-forming activities (Wang, Gotthelf & Lang 2002). In order (1) to understand the physical urrungement and interactions between the stellar and interstellar components, (2) to put constraints on distances to radio continuum features in the field of view, and (3) to make detailed images and estimates of the HI opacity toward well-studied Galactic center sources, we have carried out a complete HI absorption study of the central 200 pc of the Galaxy. This HI absorption survey represents the highest resolution and most complete study of H I absorption toward the GC, and will form the basis for the comprehensive study of the physical line-of-sight locations and interactions between interstellar features in the GC. The 15” spatial resolution and a 2.5 km s-’ velocity resolution are a vast improvement over the previous HI absorption study of Lasenby et al. ( 1 989) which included only a single pointing toward the Sgr A complex and had spatial and velocity resolutions of 50” and 10 krn S C ’ , respectively.
2 Observations & Imaging We have observed five overlapping pointings along the Galactic plane at 2 1 cm using the Very Large Array (VLA) of the National Radio Astronomy Observatory’. The VLA primary beam at 21 cm is 30’, and thus the observed region covers 100’ x 50’ (which corresponds to the central 25Ox 125 pc of the Galaxy at a distance of 8.0 kpc). These observations were made using both the DnC and CnB array configurations, giving a final spatial resolution of 15”. 127 channels were used with a total bandwidth of 1.5 MHz, corresponding to -300 km s-lof velocity coverage, centered at 0 km s-*, and a velocity resolution of 2.5 km s-’. The data were calibrated using the standard AlPS routines. The line-free channels (channels 1026) were used to fit a constant continuum level, and continuum subtraction was performed using the AlPS task UVLSF. The data were deconvolved jointly using the rniriad software package to take full advantage of the additional sensitivity from overlapping fields. Each plane of the HI absorption cube was cleaned to a uniform level using mossdi in miriad. N
N
3 Assembling an HI Absorption Catalog 3.1 Continuum Emission Figure 1 shows the continuum image constructed from -10 line-free channels. In addition, the wide field of view provided by the five mosaicked pointings allows comparison with large-scale surveys at other frequencies, such as the 90 cm image of LaRosa et al. (2000). The high sensitivity and resolution of the continuum image also allowed us to identify nine new compact sources. 3.2 HI Absorption
The continuum image has been used to guide the assembly of a catalog of continuum-weighted, line-tocontinuum HI absorption spectra towards -40 distinct continuum sources in the field. Typical rms noise for the spectra are in the range of 0.01-0.05, where the units are line-to-continuum ratio. Figure 2 shows a sample of two integrated and continuum-weighted HI profiles toward Sgr A East and Sgr A West. As Figure 2 shows, a distinct absorption feature at -40-50 km s-’ is present in the Sgr A East spectrum, but not in that of Sgr A West. This is consistent with the known interaction of Sgr A East with the so-called 40
‘
The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.
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Fig. 1 21 cm continuum image showing the cenfral 250 x 125 pc (100’ x 50‘) of the Galaxy, with a spatial resolution of 15”. The greyscale represents total intensities ranging from 0 to 125 niJy beam-’. The inset image is also a 20 cm VLA image of the central SO pc of the Galaxy (from Lang, Moms & Echevdma 1999) showing the key components in the active Radio Arc region of the GC.
km s-’cloud, and with previous conclusions that Sgr A West lies at the near cdge of Sgr A East (though still embedded in the Sgr A East shell). Figure 3 shows a sample of the continuum image (with the region used for the H I profile marked by greyscale) and corresponding HI absorption profile. Many of the profiles have very complex velocity structure. Similar profiles for -40 sources provide an overall sample o f the dataset and can guide more detailed analysis. Wc arc currently generating optical depth profiles for these regions and fitting these components to solve for the central velocities of components in each profile for identifying GC and line-of-sight features toward each of thc 40 regions.
4
C. Lang et al.: HI towards the Galactic Center
Sgr A East HI spectrum
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Sgr A West HI Spectrum
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Fig. 2 (left panel): Integrated, continuum-weighted HI absorption spectrum for Sgr A East region: (right panel): Integrated, continuum-weighted HI absorption spectrum for Sgr A West region
Sgr El1 HI Spectrum
7
Fig. 3 (left panel): Continuum image (greyscale >5u) used for determining the continuum-weighted profile of HI ahsorption toward Sgr B 1 ; (right panel): Integrated, continuum-weighted HI absorption spectrum for Sgr B1 region.
Astron. Ndchr./AN 324, No. S1 (2003)
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4 Preliminary Science Results 4.1
The Atomic Component of the Arched Filament Complex
Preliminary results from the HI absorption cube already illustrate that this dataset is rich in detailed velocity and opacity information. Figure 4 shows sample data from our HI absorption survey toward a portion of the Arched Filament HII complex. This H I I complex represents one of the most active sites in the GC, with curved filamentary arcs of ionized gas tracing out the edge of a dense molecular cloud. The extraordinary Arches cluster is responsible for the heating of this cloud (Lang, Goss & Morris 2001) but is embedded in the complex structure of filamentary molecular and ionized gas (Lang, Goss & Morris 2002). The presence of two components of atomic gas (at V N -25 kin spl and -40 km s-') and higher HI opacities toward the W1 and W2 filaments indicate that part or the atomic gas lies on the near side of this complex and that the cluster is indeed embedded within this complex. The molecular gas is presumably separated from its ionized edge by a layer of atomic material which has been photodissociated by the ionizing flux of the cluster and must have a distribution such that some of the cloud surfaces along our line of sight have not been exposed to the ionizing radiation. Figure 5 shows a sketch (from above the cloud) of a possible arrangement of various components (molecular, atomic and ionized gas in addition to the stellar cluster) that can explain the velocity profiles toward different parts of this complex (from Lang, Goss & Morris
2002).
Fig. 4 A sample of Hr absorption profiles lrom the new HI dataset toward the Arched Filament complex, which consists of molecular, ionized, atomic gas components and the dense Arches stellar cluster. The central panel shows the identification of features in 8.3 GHz radio continuum from Lang, Goss & Morris 2001) and the left and right panels show HI absorption profiles integrated over two of the filamentary features, E2 and the W1 and W2 regions. The idenfitication of the HI components are given in the figure.
C. Lang et al.: HI towards the Galactic Center
6
Ionized Gas
4 Observer
3
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Fig. 5 A schematic diagram showing a view of the ionized, molecular, atomic, and stellar components in the Arched Filament complex from a position “above” the molecular cloud, looking down its long axis. The dotted lines represent certain sightlines through the complex where complicated, double profiles are observed in the molecular and ionized gas. The numbering scheme (1-4) corresponds to sightlines of the CS(2-1) and H92cu profiles in Figure 10j, h, g, and d from Lang, Goss & Morris (2002).
4.2 Constraining Distances to Galactic Center Sources The unusual nature and morphology of the radio continuum features detected across the central 200 pc of the GC still provides the strongest case for these features being related to the center of the Galaxy. Such peculiar phenomena (e.g., the Radio Arc region comprised of the unusually-shaped “Sickle” and “Pistol” HII regions as well as the many-stranded magnetic NTFs) are not detected in radio emission toward any other region of the Galactic plane (Gray 1994). So far, we have relied on this morphological uniqueness to indicate that these sources actually reside within the central parts of the Galaxy. However, the radio continuum alone does not provide a way to adequately constrain the location of these features along our line of sight or their arrangement with respect to each other. Because of velocity crowding near .f=O degrees, the method of using galactic rotation to determine distances to GC sources can also not be used. However, the detection and identification of known HI absorption components (such as the “3 kpc-arm” and the “expanding molecular ring”; see Introduction) present in the spectra of known radio continuum sources can be used to place constraints on the distance to the radio continuum sources. Roy (2003) has recently demonstrated that distances to several of the GC magnetic filaments (e.g., Sgr C, G354.54+0.18 and G359.79+0.17) can be well-constrained to be within a few hundred pc of the GC region from measurements of HI absorption using data from the Giant Metrewave Telescope in India. We plan to carry out such anaylsis for our catalog of sources.
Astron. Nachr./AN 324, No. SI (2003)
4.3
I
A Wide HI Component Toward Sgr A
The HI absorption profiles toward parts of the SgrA complex show the presence of a large velocity dispersion (LVD) (AV -50 km s-') HI component. Several followup observations were made with the VLA in A,B,C and D array configurations using a much broader velocity coverage of -600 km s-lin order to try to characterize this unique HI component. The new data show that a low-level HI feature (~-0.3210.12), centered at -4&15 km s-lis distributed over the central I0 x 5 pc (4' by 2'). This dispersion is ten times larger than that observed in most of the cold, diffuse HI concentrations. The dispersion also appears to increase away from SgrA' indicating that the LVD HI is likely to be rotating around a distributed mass near SgrA* (Dwarakanath et al. 2003).
5
Ongoing Work on the Galactic Center HI Absorption Survey
The analysis of this rich dataset has just bcgun and the following steps outline what is planned for the thorough analysis and discussion of this H I Absorption survey: Assemble a complete catalog of HI absorption profiles toward all continuum sources in the central 200 pc. Identify the well-known sources of H I absorption toward all sources in the G C region. Such identification will also rely heavily on the molecular data of Oka et al. (1998) and other GC datasets, such as recombination line data. Constrain the distances to as many continuum sources as possible using this multi-wavelength HI absorption component identification. Construction of images showing the distrihution of HI opacity ovcr appropriate velocities. Determination of physical associations by careful comparisons of the velocity signatures of the HI , molecular and ionized gas structures. Synthesis of the HI line absorption data with the available multi-wavelength datasets is the key to determining physical interactions. 0
Quantitative estimates of the energy balance in stellar and gas components can also be derived using models of massive star energy input or Norman & Ikeuchi (1989).
Acknowledgements The authors would like to thank K. S. Dwarakanath for his assistance with putting together the catalog and analyzing the HI profiles, and S. Kim for assistance with the initial imaging. References Bally, J., Stark, A.A., Wilson, R.W., & Henkel, C. 1988, ApJ, 324, 223 Binney, J., Gerhard, O., Stark, A., Bally, J., & Uchida, K. 1991, MNRAS, 252, 210 Cohen, M. & Davies, R. 1979, MNRAS, 186,453 Dickey, J. & Lockman, F. J. 1990, Annual Reviews of Astronomy & Astrophysics, 28,215 Dwarakanath, K., Goss, W.M., Zhao, J.H. & Lang, C.C., 2003, ApJ, submitted Egan, M. P.. Shipman. R. F., Price, S. D., Carey, S. J., Clark, F. 0. and Cohen, M. 1998, ApJ, 4Y4, L199 Gray, A. 1994, MNRAS, 270, 822 Lang, C.C., Morris, M.. & Echevarria, L. 1999, ApJ, 525,727 Lang, C.C., Goss, W.M. &Moms, M. 2001. ApJ, 121, 2681 Lang, C.C., Goss. W.M. & Moms, M. 2002, AJ, 124, 2677 LaRosa, T. N., Kassim, N. E.. Lazio, T. J. W. and Hyman, S. D. 2000, AJ, 119,207 Oka, T., Hasegawa, T., Sato, F., Tsuboi, M., & Miyazaki, A. 1998, ApJS, 118,455 Norman, C. A. & Ikeuchi, S. 1989, ApJ, 345, 372 Roy, S. 2003, A&A, in press. Wang, Q. D., Gotthelf, E.V & Lang. C.C. 2002, Nature, 415, 148
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Astron. NachrJAN 324, No. S1,9- 16(2003) / DO1 10.1002/asna.200385077
High Resolution, High Sensitivity Imaging of the Galactic Center at 330 MHz. Michael E. Nard*'.*, Crystal L. B r ~ g a n * *Scott ~ , D. Hymad, T. Joseph W. Lazio', Namir E. Kassim' ,T.N. LaRosa5, K. Anantharamaiah***6,and Neboja Duric2
' Naval Research Laboratory, Code 7213, Washington, DC 203755351 USA ' University of New Mexico, Department of Physics and Astronomy, 800 Yale Blvd.
'
NE, Albuquerque, NM, 87131, USA National Radio Astronomy Observatory, 1003 Lopezville Rd. Socorro, NM, 87801, USA Department of Physics, Sweet Briar College, Sweetbriar, VA, 24595, USA Department of Biological & Physical Sciences, Kennesaw State Univ, 1000 Chastain Rd., Kennesaw, GA 30144 USA Raman Research Institute, C.V. Raman Avenue, SaddShiVanagdr Post Office, Bangalore 560-080, India
Key words Galactic Center, Low Radio Frequencies, Sgr A*, Non-Thermal Sources, Wide-field Imaging.
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Abstract. We present results derived from a wide field, sub-arcminute resolution VLA image of the Galactic Center region at 330 MHz (A = 90 cm). With a resolution of 7" x 12" and an rms noise of 1.6 mJy bean-', this image represents a significant increase in resolution and sensitivity over the previously published VLA image at this frequency (eg. LaRosa et al. 2000). The improvement in sensitivity has significantly increased the census of sinall diameter sources in the region, resulted in the detection of two new Non-Thermal Filaments (NTFs) and 18 new NTF candidates, and resulted in the lowest frequency (tentative) detection of Sgr A*.
1 Introduction The Galactic Center (GC) was first imaged with the VLA at 330 MHz in 1989 (Pedlar et al. 1989; Anantharamaiah et al. 1991) and represented a major improvement in sensitivity and resolution over past meter wavelength images. However, wide field imaging algorithms at the time were unable to compensate for the non-coplanar nature of the VLA. Hence the full primary beam of the VLA at 330 MHz (-- 2.2' radius) was not imaged and only the very center of the GC region was studied. The G C image published by LaRosa et al (2000) represented a major improvement. It led to the discovery of many new sources, and provided an unparalleled census of both extended and small diameter, thermal and non-thermal sources within 100 pc (projection) of the GC. This major step forward was afforded by significant advances in wide-field imaging algorithms, coupled with greatly increased computational power. However that effort fell short of utilizing the full resolving power of the VLA and the commensurate greater sensitivity that it would have afforded. Since those data were obtained and published, significant improvements in software, hardware, and computational power have continued 10 be realized. This inspired us to revisit the G C at 330 MHz.
* Corresponding author: e-mail: Michael.NordOnrl.navy.mil,Phone: +1 202767 7310, Fax: +01202404 8894 * * NRAO Jansky Postdoctoral Fellow. * * * Deceased.
@ 2003 WILEY-VCH Verlag GmhH & Co K O I A Weinhem
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M. E. Nord et al.: High Resolution, High Sensitivitv Imagine of the Galactic Center at 330 MHz.
Fig. 1 The Galactic Center Region at 330 MHz as imaged with the VLA. This image has aresolution of 7" x 12" and an rms noise of 1.6 mJy bean-', and is a sub-image of a much larger image. This sub-image is roughly 0.8" x 1.0". The entire image covers an area of nearly 15.5 square degrees. Two new Non-Thermal Filaments and 18 new NonThermal Filament candidates were discovered in this image. A full explaination of how the image was constructed as well as the full details on results will be published in Nord et al., in prep.
Astron. Nachr./AN 324. No. S1 (2003)
I1
We present here the image generated from ncw A and B configuration data sets, which are appropriate for generating an image with a minimum of confusion noise and maximum sensitivity to relatively smaller scale (< 1’) structure. For the first time, the entire GC region contained by the primary beam of the VLA has been imaged at the maximum possible resolution and sensitivity. The image has a resolution of 7” x 12’’ and an rms noise of 1.6 mJy beam-’, an improvement by roughly a factor of 5 in both sensitivity and resolution over the LaRosa et al. (2000) image. Figure 1 is a subimage of the central GC region. Here we describe new results pertaining 10 Sgr A* and to newly discovered NTFs and NTF candidates. A full description of data reduction techniques and results pertaining to compact source distribution, spectral indeces, greater detail in extcnded sources, and Galactic Center scattering will be presented in Nord et al. (2003, in preparation). Efforts to image the GC from the combination of data obtained in all four VLA configurations, which is more appropriate for the study of more extended (> 10’) emission, are under way.
2 Observations Three sets of observations were obtained. Thc tirst was observed at 330 MHz in the A configuration of the VLA (maximum baselines 35 km) in Octoher of 1996. Data were obtained in two IF’Swith dual circular polarization, a bandwidth of 3 MHz, and the total bandwidth was split into 32 channels in order to enable later radio frequency interference (RFI) excision as well as to mitigate the effects of bandwidth smearing. The 1996 data were obtained from a series of observations designed to find candidate GC pulsars (Lazio & Cordes, in prep.). Between March 1998 and May 1999, the Galactic Center was observed in all four VLA configurations at a single IF centered at 330 MHz with 3 MHz bandwidth and 32 channels, while the other IF was dedicated to 74 MHz. All observations wcre full synthesis which tracked the Galactic Center as long as it was visible to the VLA (- 6 hours). and unlike the archival data re-processed by LaRosa et al. (2000). all these new data were obtained using all 27 antennas of the VLA. These details are summarized i n Table 1. Images made from the 74 MHz data are presented in Brogan et al., these proceedings. ~4
Table 1 Observational Summary
Epoch 1996 October 1998 September 1998 March
VLA Configuration A
B A
I/
(MHz) 332.5 327.5 327.5
3 SgrA”
#of IFs 2 1 1
Integration (Hours)
Beam
5.55 5.47 5.38
9” x 5’’
36” x 20” 9’l x 5”
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Below 10 GHz, Sgr A* has a nearly flat spcctrum with a flux density of 0.6 Jy. In 1989, Pedlar et al. suggested that at 330 MHz, the flux density of Sgr A* could be no more than 100 mJy. It is not fully understood why the flux density at 330 MHz should be so low. Pedlar et al. suggest that this is due to foreground free-free absorption of optical depth T 2.5, while Beckert et al. (1996) suggest that the turn over may be modeled by synchrotron self-absorption intrinsic to the source itself. In order to more accurately measure the flux density of Sgr A*, confusion with the large scale components of flux i n the region must be considcrcd. The data presented here are already chosen to minimize the contribution of large scale flux. The super-uniform weighting scheme (Briggs et al. 1999) can further reduce the contribution of short spatial frequencies in the image, thereby decreasing the size of the synthesized beam at the cost of increasing the overall noise level in the image. The region around Sgr A* was imaged with super-uniform weighting and through this technique the size of the synthesized beam was 10 mJy beam-’. decreased to 3.4“ x 8.2” (RA, DEC) while increasing the rms noise of the image to The super-uniform image is shown in Figure 2.
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M. E. Nord et al.: High Resolution, High Sensitivity Imaging of the Galactic Center at 330 MHz.
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01
Fig. 2 Super-uniformimage of the Sgr A* region at 330 MHz. Image display is inverted: white areas represent areas of low flux due to absorption. The grey scale is linear between -2 and 100 mJy bean-'. The resolution is 3.4" x 8.2'' and the rms noise is -10 mJy beam-'. Note the area of diffuse eniission at
Fig. 3 Model created from 6 cm data of the Sg A* region as it should appear at 330 MHz. Image display is inverted; white areas represent areas of low flux due to absorption. See Section 3 for details.
the center. The oval represents the location, predicted scattering diameter and orientation of Sgr A* at 330 MHz.
The scatter-broadened size of Sgr A* at 330 MHz was estimated using the measurements at 20.7 cm by Yusef-Zadeh et al. (1994). Taking the 20.7 cm size of 624 x 350 milliarcseconds and scaling by A', the size of Sgr A* is calculated to be 11.8" x 6.6". The position angle was assumed to be constant at -80". The scatter-broadened size, location, and position angle of Sgr A" are represented by the oval in Figure 2 Of note in the super-uniform image is a diffuse region of flux having a peak intensity within 5" of the location of Sgr A*. Slices along the major and minor axes of the scattering disk are shown in Figures 4 and 5. These slices strongly suggest that there is a narrow region of 330 MHz flux near the position of Sgr A*. A gaussian fit with baseline subtraction to this region gives a source size of 19" f 4 x 11" f 1, a position angle of -30", a peak intensity of 50110 mJy beam-' and a flux density of 330f100 mJy. The expected size of Sgr A' convolved with the super-uniform beam gives an expected source size of 10.5'' x 12". With the correct position, nearly the correct size, and a flux density near that of higher frequency values, we decided to test the hypothesis that this diffuse source is Sgr A*. In order to check that the source being observed is in fact Sgr A*, a 330 MHz model of the Sgr A* region was created from a 6 cm (v = 5 GHz) image. Starting with a 6 cm image (Yusef-Zadeh et al. 1987), a 0.75 Jy model of Sgr A* was subtracted. The remaining flux should be ionized gas from the 'thermal spiral', Sgr A West (Lo & Claussen 1983). As this ionized gas will be in absorption against the non-thermal flux from the Sgr A East supernova remnant at 330 MHz, the 6 cm image is inverted, so that positive flux becomes negative. A model of Sgr A* with a flux density of 0.75 Jy, and the appropriate 330 MHz scatter-broadened size was added back into the image, and then the image was convolved to have the resolution of the super uniform image at 330 MHz. The resulting image is shown in Figure 3. This model shows significant similarities to the observed super uniform image. The region of absorption caused by the thermal gas matches nearly exactly. More interestingly, at thc position of Sgr A*, there is a region of extended flux similar in size, shape and location as in the 330 MHz super uniform image. Morphologically, N
Astron. Nachr./AN 324, No. S1 (2003)
Fig. 4 Slice through the assumed major axis of Sgr A* in the super-uniform image (Figure 2). The slice is centered at 17h 45"' 40" -29"00'28"(52000) with a position angle of -90".
13
Fig. 5 Slice through the assumed minor axis of Sgr A* in the super-uniform image (Figure 2). The slice is centered at 17h 45'" 40" -29"00'28"(52000).
the diffuse region in the 330 MHz super uniform image at the position of Sgr A* appears similar to what this model predicts. Before identifying this diffuse source as Sgr A*, the possibility that the source is background flux from the Sgr A supernova remnant was investigated. Figure 6 is a greyscale image of the Sgr A* region with the 6 cm, Sgr A* subtracted image of the previous paragraph superimposed in contours, and the position, scatter broadened size, and position angle of Sgr A* represented by an oval. At the position of Sgr A", the 6 cm intensity is near maximum, suggesting that this should be the area of maximum absorption at 330 MHz. However, what is seen at 330 MHz is a region of extended emission. This observation makes it unlikely that the flux originates from behind the absorbing ionized region. We therefore conclude that we are detecting Sgr A* at 330 MHz for the first time. The position of maximum intensity, size, and shape of the dirfuse region all agree with our model of how the source should appear at 330 MHz. We fit a maximum intensity of 504110 mJy beam-', and a flux density of 3301100 mJy. However, we leave open the possibility that the flux is from, or is contaminated by, background non-thermal flux from the Sgr A East supernova remnant. See Figure 7 for a radio (0.33 < v < 2 3 GHz) spectrum of Sgr A*.
4 Non-Thermal Filaments Among the most fascinating of the unique structures in the Galactic Center are the non-thermal filaments. NTFs are remarkably coherent magnetic structures that extend tens of parsecs and maintain widths of only a few tenths of parsecs (e.g., Lang et al. 1999). It has been hypothesized that the NTFs are part of a globally ordered, space filling magnetic field (e.g., Morris & Serabyn 1996). If so, they would be the primary diagnostic of the GC magnetic field. An alternative idea is that the NTFs are magnetic wakes formed from the amplification or a weak global field through a molecular cloud galactic center wind interaction (Shore & LaRosa 1999; see also LaRosa et al., these proceedings).
M. E. Nord et al.: High Resolution, High Sensitivity Imaging of the Galactic Center at 330 MHz.
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Fig. 6 Greyscale image of the Sgr A* region at 330 MHz. Image display is inverted white areas represent areas of low flux due to absorption. The grey scale is linear between -5 and 400 mJy beam-’. The resolution is 6.81” x 12.58” and the rms noise is ~ 1 . mJy 6 beam-’. Contours are 6 cm with a resolution 3.39” x 2.93” and levels of 1, 2, 3, and 3 times SO mJy beam-’. A model of Sgr A’containing 750 mJy has been subtracted from the 6 cm contours. Note the area of diffuse emission at the center. The oval represents the location, predicted scattering diameter and orientation of Sgr A*at 330 MHz.
Fig. 7 The low frequency (0.33 < u < 23 GHz) spectrum of Sgr A*. The 0.61 GHz value is from Subhashis Roy (private communication). The 1.4 < u < 23 GHz values are from Zhao et al. (2001). Also included is the projected LOFAR detection limit at 240 MHz (see Kassim et al., these proceedings).
Nine NTFs were known before this work was completed. Of those nine, we detect eight as we do not have sufficient surface brightness sensitivity to detect G359.85+0.39 (LaRosa et al. 2000, 2001). In addition to higher resolution and higher sensitivity measurements of the known NTFs, we report the discovery of two new NTFs and 18 new NTF candidates. For the two new NTFs, confirmation of NTF status was made by observations of 6 cm polarization which is reported on in greater detail in LaRosa et al. (2003a, these proceedings). If all are confirmed, these represent a tripling of the number of known NTFs. Table 2 summarizes the new NTFs and NTF candidates. Of particular interest is the orientation with respect to the Galactic Plane of these new NTF candiates, and what this can tell us about the GC magnetic field. This topic as well as images of several of the new NTFs and NTF candidates is covered In LaRosa et al., (2003a, thcse proceedings). Figure 8 is a histogram of the NTF’s maximum intensity. We use intensity instead of flux density for two reasons. Firstly, detection of these sources is based on maximum intensity, not overall flux and secondly baseline subtraction can be very difficult for very extended sources that pass through regions of diffuse flux, so the total integrated flux density of the NTFs is uncertain. Though there are not enough sources to definitively fit a surface brightness curve, the increase towards low surface brightness rises faster than SP2. By increasing the sensitivity by a factor of -5, we have tripled the linear and appears to be number of known NTFs, suggesting that the total number of NTFs rises at minimum as N ( S ) S-3/5.
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N
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Astron. Nachr./AN 324, No. S I (2003) Table 2 New Non-Thermal Filaments and Candidates
Name
Intensity (mJy beam
NTF 359.12+0.66" NTF 359.22-0.16' NTF 359.33-0.42 NTF 359.36+0.09 NTF 359.40-0.03 NTF 359.40-0.07 NTF 359.43+0.13 NTF 359.59-0.34 NTF 359.66-0.1 1 NTF 359.85-0.02 NTF 359.86-0.24 NTF 359.88-0.07 NTF 359.90+0.19 NTF 359.99-0.54 NTF 0.02+0.04 NTF 0.06-0.07 NTF 0.37-0.07 NTF 0.39+0.05 NTF 0.39-0.12' NTF 0.43+0.01
11.9 23.4 13.9 10.7 11.9 40.6 18.8 20.8 9.9 8.5 11.2 33.8 11.9 9.4 22.3 10.5 14.1 17.8 16.1 11.6
')
Flux Density (dy) 647.8 268.5 80.6 65.2 93.8 229. I 264.5 188.2 226.0 172.5 205.1 930.0 129.2 88.8 227.8 162.6 128.1 231.5 731.2 43.9
Size
Plane Angle"
(9
(")
15.6 x 0.2 1.8 x 0.5 2.0 x 0.2 2.1 x 0.2 1.6 x 0.2 1.7 x 0.3 2.4 x 0.3 2.3 x 0.2 3.5 x 0.5 1.8 x 0.2 8.1 x 0.2 1.6 x 0.2 2.4 x 0.2 8.6 x 0.2 2.0 x 0.3 2.1 x 0.2 1.1 x 0.3 4.1 x 0.3 10.1 x 0.3 1.6 x 0.3
35 55 55 60 5 40 0,90d 25 20 90 35 5 35 30 0 15 5 5 5 5
Plane Angle is the angle of the NTF with respect to the normal to the Galactic Plane. This filament was first detected at higher frequencies ( M . Morris 2002, private communication). Source observed to have significant 6 cm polarization, and therefore is confirmed an NTF (LaRosa et al 200%. these proceedings). This source may be two interacting NTFs with orientations of Oo and 90' to the Galactic plane.
We conclude that just the tip of the NTF luminosity distribution is being detected and we hypothesize that there may be hundreds more NTFs in the G C region (see Kassim et al. 2003, these proceedings).
5
Conclusions
Rcsults from a new high resolution, high sensitivity imaging of the Galactic Center region at 330 MHz were presented. In addition to more than tripling the number of NTF candidates, we report the tentative lowest frequency detection of Sgr A*. Further results will be presented in Nord et al., in prep. In addition to shorter baseline VLA data that is LO be added to this dataset, time has been allocated on thc Green Bank Telescope for adding single dish information to these data. This will allow us to make the connection from the structures seen in our interferometer images to those larger scale features seen on single-dish images which have been interpreted as being related to, for example, episodic periods of star-burst activity that can also be traced in X-ray images (Sofue 2000).
Acknowledgements This project was originally conceived under the guidance of K. Anantharamaiah. Anantha passed away during the data reduction phase of this project and will be missed greatly. The authors would like to thank Manana S . Lazarova and Jennifer L. Neureuther, students at Sweet Briar College for their assistance in small diameter source location and quantification.
M. E. Nord et al.: High Resolution, High Sensitivity Imaging of the Galactic Center at 330 MHz.
16
Fig. 8 NTF maximum intensity histogram, including the 9 previously known NTFs, 2 newly discovered NTFs and 18 NTF candidates.
Basic research in radio astronomy at the NRL is supported by the Office of Naval Research. S.D.H was supported by a grant from the Jeffres Memorial Trust and Research Corporation. The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.
References Anantharamaiah, K. R., Pedlar, A,, Ekers, R. D., & Goss, W. M. 1991, MNRAS,249,262 Beckert, T., Duschl, W. J., Mezger, P. G., & Zylka, R. 1996,A&A,307,450B Briggs, D. S., Schwab, F. R., & Sramek, R. A. 1999, ASP Conf. Ser. 180: Synthesis Imaging in Radio Astronomy 11, 127 Brogan et al. 2003, these proceedings Kassim et al. 2003, these proceedings LaRosa, T. N., Kassim, N. E., Lazio, T. J. W., & Hyman, S. D. 2000, AJ, 119, 207 LaRosa, T. N., Lazio, T. J. W., & Kassim, N. E. 2001, ApJ, 563, 163 LaRosa et al. 2003a, these proceedings LaRosa et al. 2003b (in preparation) Lazio & Cordes 2003 (in preparation) Lo, K. Y. & Claussen, M. J. 1983, Nature, 306,647 Liszt, H. S. & Spiker, R. W. 1995, ApJS, 98,259 Moms, M. & Serabyn, E. 1996, ARA&A, 34,645 Nord et al. 2003, (in preparation) Pedlar, A,, Anantharamaiah, K. R., Ekers, R. D., Goss, W. M., van Gorkom, J. H., Schwarz, U. J., & Zhao, J. 1989, ApJ, 342,769 Shore, S. N. & Larosa, T. N. 1999, ApJ, 521,587 Sofue, Y. 2000, ApJ, 540, 224 Yusef-Zadeb, F. & Morris, M. 1987, ApJ, 320,545 Yusef-Zadeh, E, Cotton, W., Wardle, M., Melia, F., & Roberts, D. A. 1994, ApJ, 434, L63 Zhao, J., Bower, G. C., & Goss, W. M. 2001, ApJ, 547, L29
Astron. Nachr./AN 324, No. S 1, 17 - 24 (2003)/ DO1 10. I002/asna.2003F5023
Spatially Resolved Very Large Array 74 MHz Observations Toward the Galactic Center C. L. Brogan*’,M. Nord’, N. Kassim’, J. Lazio’, and K. Anantharamaiah**3
’ National Radio Astronomy Observatory, 1003 Lopezville Rd., Socorro, NM 87801 Navel Research Laboratory,4555 Overlook Ave., Code 7213, Washington, DC, 20375 ’ Raman Research Institute, C. V. Raman Avenue, SadashivanagarPost Office, Bangalore 560-080, India
Key words Cosmic Rays, Galactic Center, H ‘ 1 regions, Supernova Remnants PACS 04A25
We present the highest resolution and sensitivity low frequency image (< 300 MHz) of the Galactic center to date using the Very Large Array at 74 MHL in its A, B, C, & D configurations. The resulting images have a resolution of 2.1’ x 1.2 and a dynamic range of 400. From this data we have been able to identify a region of enhanced 74 MHI emission about 5” in extent that is coincident with the high density molecular gas surrounding the Galactic ccntcr known as the Central Molecular Zone. In addition to giving an unprecedented view of the extended nontherinal emission surrounding the Galactic center, the 74 MHL image shows deep free-free absorption across the Galactic center itself, as well as, part of the Galactic center radio lobe, and a number of H II regions in the field. This absorption is due to ionized thermal gas in front of, or in some cases embedded in, the nonthermal Galactic center (CC) emission. Such absorption allows us to unambiguously place some of the H I I regions in the direction of the GC along the line of sight for the first time. The morphology, nature, and relationship to the Galactic center of the 74 MHz absorption and emission is discussed. N
1 Introduction Low frequency (< 100 MHz) observations of the Galactic center (GC) provide a unique window on distinct synchrotron sources like supernova remnants, as well as, the diffuse G C synchrotron emission. In addition, because ionized thermal gas can absorb low frequency synchrotron emission, these data can also be used to unambiguously place H I I regions along the line of sight toward the GC. In the past, radio observations bclow I00 MHz were of limited usefulness due to their low resolution and sensitivity. The recent addition to the Very Large Array (VLA) of a 74 MHz receiver system has allowed for high resolution and sensitivity low frequency observing for the first time. The 25 mcter size of the VLA antennas also provides a very large 14” field of view at 7 4 MHz, thus allowing for a complete picture of the whole GC region. Since the detection of absorption in radio continuum data is relatively rare, we briefly describe the physics behind such observations. For objects observed in 74 MHz absorption with an interferometer, the where , T, is the electron tempcraturc of the observed brightness temperaturc will be T o h s = T, - T G , ~ ionized thermal gas (H I I region) and T c : ,is~ the temperature of the Galactic synchrotron crnission behind the H 11 region. Therefore, if T G , 6 x
Astron. Nachr./AN 324, No. S 1 (2003)
35
Table 1 Galactic X-ray Point Sources
log( L y )" Spectru m0 lOg(C'Kg . - ' ) MS Stars' 25 - 30 :I AT < 1 keV Plasma 29 31.1 YSOS kT = 1 - 10 keV Plasma RS CVdAlgol 2!) :31 7 kT = 0.1 2 keV Pla5ma 31 kT = 0.1 - 6 keV Plasma WR/O Stars 20.5 - 32 (1 k T = 1 25 keV Plasma cvs Pulsars 29 3 - :K) r = 1 2.5 PL ; kT = 0 . 3 keV BB NS LMXBs 31.6 - 38 kT 0 3 keV BB ; r = 1 2 PL'' 30 - : 90% confidence in three energy bands above 2 keV. along with the colors expected for an absorbed power-law spectrum. Simulations indicate that the thermal models expected for the coronal sources and for most CVs all produce hard colors less than 0.1. This includes ionized plasmas with k T < 25 keV, blackbodies with kT < 2 keV, or Bremsstrahlung emission with kT < 50 keV (see Table 1). Although these softer sourccs should be the most numerous in our field, only a small fraction are expected to be bright enough above 2 keV to observe through the Galactic absorbing column of 6 x 10" cm-' of H. On the other hand, over half of the Galactic Center sources cluster i n a region consistent with absorption columns l o g ( N ~> ) 22.5, and very flat spectra with photon indices r < 1 (where negative values indicate rising numbers of photons with energy). Such hard spectra are unusual for X-ray point sources. but have been seen previously from cataclysmic variables (CVs) containing magnetized white dwarfs (polars and intermediate polars; e.g., Ezuka & Ishida 1999) and from magnetized neutron stars accreting from the winds of massive companions (High-Mass X-ray Binary [HMXB] pulsars; e.g., Campana et al. 2001). Supporting this hypothesis, 8 of the brightest 285 sources exhibit coherent X-ray pulsations with periods ranging from 300 s to 4.5 hours (Muno et al. 2003b, in preparation).
3 Spatial Distribution and the Star Formation History Within about 8' from the aim point, we estimate that we can detect all sources with photon fluxes greater than 5 x l o p 7 photons cmp2 s-* with a signal-to-noise of at least 71, 3 i n the 2.0-8.0 keV band. About 40%' of the Galactic Center sources detected within 8' of Sgr A* have photon fluxes greater than this value. In Figure 3 we plot the number of Galactic Center sources above this flux limit per unit solid angle as a 1
M. P. Muno et al. Galactic Center Chandru Sources
36
-0.5L 10 -1.0
L
L
I
m
-0.5
,
I
I
+ I
,
,
1
,
0.0
I
0.5
.
I
,
,
1 .0
Medium Color
Fig. 2 Comparison of the observed hard and medium colors to those expected from an absorbed power-law spectrum for sources detected with greater thdn 90% confidence in all three energy bands above 2 keV. The colors are defined as the fractional difference between the count rates in two energy bands, ( h s ) / ( h + s), where h and s are the numbers of counts in the higher and lower energy bands, respectively. The medium color is defined using counts with energies between 3.34.7 keV and 2.0-3.3 keV, and the hard color using counts between 4.7-8.0 keV and 3.34.7 keV. Open circles denote data from 39 foreground sources and filled circles from 785 sources at the Galactic Center. The crosses connected with solid lines indicate the expected colors for absorbed power laws, with the photon indices and absorption columns indicated on the plot. The median uncertainty for these sources is displayed at the bottom of the plot. We note that the sources in the upper-left comer of this Figure all have uncertainties a factor of 2-3 larger. ~
function of angular separation from Sgr A*. We have fit this surface density distribution with a power law of the form
where H is the angular separation from Sgr A* in arcminutes. Both the normalization and the power-law of 26 for 30 degrees of freedom. slope were allowed to vary. The resulting fit is acceptable, with a If we assume that these sources are distributed with spherical symmetry about the Galactic Center, the implied spatial density falls off with radius as R-2. The spatial density of stellar sources observed in the infrared also decreases approximately as R-a.o*o.3within 30 pc of Sgr A* (Serabyn & Morris 1996). This suggests that the X-ray sources lie primarily in the Nuclear Stellar Clustcr (Mezger, Duschl, & Zylka 1996; Launhardt et a1.2002) and that their spatial distribution traces that of infrared stars. This suggests that the stellar X-ray sources can be used to understand the star formation history in the inner tens of parsecs of the Galaxy, where it is uncertain how the large tidal forces and the milliGauss magnetic fields affect star formation, and where many traditional observational tracers of star formation have been difficult to find (Mezger et al. 1996; Morris 1993; Serabyn & Morris 1996). For instance, if there are magnetic CVs among the hard X-ray sourccs, they would trace low-mass stars in the nuclear bulge (Warner 1995). On the other hand, if some of them are wind-accrcting neutron stars, they would provide an important constraint on the amount of massive stars formed near the Galactic Center in the last lo7 - loy years (Pfahl, Rappaport, & Podsiadlowski 2002).
xz
Astron. Nachr./AN 324. No. S 1 (2003)
60
"
'
I
"
"
'
'
'
676 sources with F
in
OF
I
0
,
I
1
,
2
,
,
>
50x10.'
,
.
~~
I
4
37
1
"
'
ph cm-' s-'
, +
6
a
Offset (arcmin)
Fig. 3 Surface density of Galactic Center poinl sources as a function of offset angle (8) from Sgr A*. The number of sources in each annulus was divided by the solid angle over which a source could he detected above 5 x lor7 photons cmr2 srl with a signal-to-noise of 3 in that annulus. The dotted linc indicates the best-fit power law, 2: cx Or' "*".').
3.1
Fig. 4 Cumulative log(N) - Iog(S) distribution of sources at the Galactic Center (filled circles,
2-8 keV) and in the foreground (open circles, 0.5-8 keV). The best-fit models determined by a maximum-likelihood method are over-plotted with solid lines. The expected extra-galactic contribution from Brandt et al. (2001) is indicated with the dashed line (see also Rosati et al. 2002).
Point Source Contribution to the Diffuse X-ray Emission
A hot (lo8 K) component of the ISM is thought to be responsible for the He-like Fe 6.7 keV emission that is observed all along the Galactic ridge. However, the temperature of this putative diffuse plasma is much higher than that typically produced in supernova shocks, and it is too high for the plasma to he gravitationally hound to the Galactic disk (Worrall et al. 1Y82; Koyama et al. 1986). If the plasma is unbound, the power required to sustain this hard Galactic ridge emission is approximately 10" erg spl, equivalent to the kinetic energy of one supernova occurring every 30 ycars (e.g., Valinia & Marshall 1998). This input would have to be provided by exotic processes, such as interactions between lowenergy (10 keV) cosmic-rays and the ISM (Valinia & Marshall 1998; Tanaka, Miyaji, & Hasinger 1999) or magnetic reconnection driven by turbulence in the ISM (Tanuma et al. 1999). However, there are currently no independent means of observing either low-cncrgy (10 keV) cosmic rays or magnetic reconnection (but sce Serabyn & Morris 1994 for the latter), so it is important to estimate the energy in the hot plasma as accurately as possible. The point sources in Figure 1 could contribute significantly to the Galactic Ridge X-ray emission, and t h u s lessen the energetic requirements on thc hot plasma. In Figure 4, we plot the cumulative number counts as a function o f flux for sources at the Galactic Center (filled circles, fluxes between 2-8 keV) and in the roreground (open circles, 0.5-8 keV), normalized to the solid angle of the survey 12 in units of arcmin-2. We have been conservative in our source selection to avoid incompleteness in our sample caused by the varying sensitivity over our image (see Muno et al. 2003 for complete details). Note that only about one-third of hoth the foreground and Galactic Center sources detected in our image satisfy our selection criteria, since thc high background in the image adds significant uncertainty to our flux measurements We focus on the distribution of Galactic Center sources, since these are most relevant to understanding the origin of the hard component of the diffuse X-ray emission. Using the un-binned flux valucs, we modeled the log AT log S distributions using the maximum likelihood technique described in Murdoch, Crawford. & Jauncey (1975). The distribution of Galactic Center sources were consistent with two power ~
M. P. Muno et al. Galactic Center Chandrcl Sources
38
laws of the form
where S is in units of photons cm-' s-', and the normalization is in units of sources arcinin-'. We estimate the total flux produced by these point sources by integrating the flux convolvcd with Equation 2, converting the photon fluxes into energy fluxes by assuming l photons cmp2 s-l = 8 x lo-" erg cm-2 s-l (2.0-8 keV). We find that point sources with fluxes greater than 3 x 10W" erg cm-' s-' contribute a mean surface brightness of 4 x erg cm-' 5-l arcmin-' over the inner 9' around Sgr A*. This is about 10%)of the diffuse flux from the inner regions of the Galaxy derived by Koyama et al. (1996; erg cmP2 s-* arcmin-') and by Sidoli & Mereghetti (2001; erg cm-2 s-' in a I " field, or 3 x 1x erg cm-' s-' in a 190 arcmin-' field, or 5 x lo-'" erg cm-2 s-' arcmin'). A similar result was obtained by Ebisawa et al. (2001) in a rcgion at I = 28" and h = 0.2". However, the steep slope of the flux distribution ( S - ' 7 , below 6 x lo-'' erg cm-2 s-', implies that lhe integrated flux from point sources in the field will diverge if the distribution extends to arbitrarily low fluxes (see Figure 4). Point sources would account for all of the diffuse emission reported by Koyama et al. (1996) and Sidoli & Mereghetti (1999) if the distribution in Equation 2 extends a factor of 40-100 lower in flux. However, from the image in Figure 1, it is clear that filamentary features contribute a significant liaction of the diffuse emission, which implies that the flux distribution in the 2-8 keV band (where most erg cm-' of the diffuse emission is observed) must turn over between fluxes of 3 x 10-'7 and 3 x s-l, or luminosities o f 2 x loz9 to 2 x 10"' erg s-' at the Galactic Center. We have also compared the spectra of the point sources and the diffuse emission, in order to constrain the relative contribution of each to the 6.7 keV iron emission from the Galactic Center. The equivalent widths of the apparently diffuse 6.7 keV emission range from 350-500 eV, even in regions that are otherwise relatively free from filamentary features (see Park et al. 2003, in these proceedings). The iron emission in the combined spectrum of the Galactic Center point sources is at the low end of this range, about 330 eV. Thus, it appears that unresolved point sources are unlikely to account for all of the 10 keV component of the diffuse Galactic X-ray emission (Muno et al. 2003c, in preparation).
4
Conclusions
We have detected 2357 X-ray point sources during 620 ks of Chundru observations of the 17' x 17' field around Sgr A* (Figure I ) . The completeness limit of our survey at the Galactic Center is about 3 x erg cm12 ~ ~ ' ( 2 -keV), 8 while sources are detected with fluxes nearly an order of magnitude lower. Only 20-100 of these sources are expected to be background AGN. The large number of sources in this field probably results from the high stellar density at the Galactic Center. Indeed, we have demonstrated that the surface density of Galactic Center X-ray sources decreases as N l/O away from Sgr A* (Figure 3). just as the surface density of infrared stars does. We have also shown that the l o g ( N )- log(S) distribution of the Galactic Center sources is very steep, rising as S-1.7near our completeness limit (Figure 4). This indicates that unresolved point sources can contribute significantly to the diffuse component of the Galactic X-ray emission. More than half of the sources for which we have spectral information are very hard, with spectra that are consistent with I' < 1 power laws (Figure 2). Such hard spectra have only been observed previously from magnetically accreting white dwarfs and wind-accreting neutron stars. If they are magnetic CVs among these X-ray sources, they would he the first low-mass stars identified in the nuclear bulge. If they are wind-accreting neutron stars, these systems would provide an important constraint on the amount of star formation that has taken place near the Galactic Center in the last lo7 - lo8 years. This highlights the importance of identifying the nature of the Galactic Center sources with more certainty. The X-ray spectral
Astron. Nachr./AN 324, No. S1 (2003)
39
and timing properties of these sources will be reported in detail in the near future, and w e are in the process of identifying these sources at radio and infrared wavelengths.
Acknowledgements We gratefully thank D. Schwartz, P. S h e , and the Churidru Mission Planning group for their efforts in scheduling thc observations in late May so that they would have nearly identical roll angles and aim points. We also thank K. Getman and F. Batter for developing software that aided us in collating and visualizing these results, V. te Velde for helping to construct a map of the diffuse background, P. Schechter for valuable advice on how to treat the log(N) - log(S) distribution, and E. Pfahl and J. Sokoloski for helpful discussions about the possible natures of these sources. This work has been supported by NASA grants NAS 8-39073 and NAS 8-00128. W.N.B. alho acknowledges the LTSA grant NAG 5-81 07 and the Alfrcd P. Sloan foundation.
References Baganoff, F. K.. Bautz, M. W., Brandt, W. N.. Chartas, G., Feigclson, E. F., Garmire, G. P., Maeda, Y., Morris, M., Ricker. G. R., Townsley, L. K., &Walter, F. 2001. Nature, 413, 45 Baganoff, F. K. et al. 2003, ApJ, 591. 891 Brandt, W. N. et al. 2001, AJ, 122, 2810 Campana, S., Gastaldello, F., Stella, L., Israel, G. L., Colpi, M., Pizzolato. F., Orlandini, M., & Dal Fiume. D. 2001, ApJ, 561.924 Ebkawa, K., Maedd, Y., Kaneda, H.,& Yamauchi. S. 2001 b, Science, 293, I633 Ezuka, H. & Ishida, M. 1999, ApJS, 120, 277 Fabbiano, G. 1989, ARA&A, 27, 87 Feigelson, E. D., Broos, P., Gaffney, J. A. 111, Garinire, G . , Hillenbrand. L. A,, Pravdo, S. H., Townsley, L., & Tsuboi, Y. 2002, ApJ, 574, 258 Giistcn, R., Walmsley, C. M., & Pauls, T. 19x1, A&A, 103, 197 Koyama, K., Maeda, Y., Sonobe, T., Takeshima, T., Taneka, Y., & Yamauchi, S. 1996, PASJ, 48. 249 Koyaina, K., Makishinia, K., Tanaka, Y., & Tsunemi, H. 1986, PASJ, 38, 121 Launhardt, R., Zylka, R., & MeLger. P. G. 2002, A&A, 384, I12 rf Maeda, Y. et al. 2002, ApJ, 570, 67 I Mewe, R., Lemen, J. R., & van den Oord, G. H. J . 1986, A&AS. 65.5 I I Zylka, R., 1996, AARev, 7. 289 Muno, M. P., Baganoff, F. K.. Bautz. M. W.. Ricker, G. R., Monis, M., Garinire, G. P., Feigeison, E. D., Brandt, W. N., Townsely, L. K., & Broos, P. S. 2003a, ApJ, 589, 225 Murdoch, H. S., Crawford, D. F., & Jauncey, D. 1973. ApJ, 183, 1 Pavlinsky, M. N., Grebenev, S. A.. & Sunyaev, R. A. 1994, ApJ, 425, 110 Pfahl, E., Rappaport, S., & Podsiadlowski, P. 2002, ApJ, 571, L37 Predehl, P. & Truemper, J. 1994, A&A, 290, L29 Raymond, J. C. & Smith, B. W. 1977, ApJS, 35,419 Rosati, P. et al. 2002, ApJ, 566, 667 Sakano, M., Koyama, K.. Murakami, H., Maeda. Y.. & Yamauchi, S. 2002, ApJS, 138, 19 Serabyn, E. & Morris, M. 1994, ApJ, 424, L91 Serahyn, E. & Morris, M. 1996, Nature, 382,602 Sidoli, L., Belloni, T., & Mereghetti, S . 2001. A&A, 368, 835 Sidoli, L. & Mereghetti, S. 1999, A&A, 34’). L49 Sidoli, L., Mereghetti, S., Israel, G. L., Chiappctti, L., Treves, A., & Orlandini, M. 1999, ApJ, 525, 215 Tanuma, S., Yokoyama, T., Kudoh, T., Matsumoto. R., Shihata, K., & Makishima, K. 1999, PASJ, 51, I61 Tanaka, Y., Miyaji, T., & Hasinger, G. 1999. Astron. Nachr., 320, 181 Valinia, A. & Marshall, F. E. 1998, ApJ, 505, 134 Wdmer, B. 1995, Curaclv,rmic Variable .Star.\, Cambridge University Press Watson. M. G., Willingale, R., Hertz, P., & Grindlay, J . E. 1981, ApJ, 250, 142 Worrall, D. M., Marshall, F. E., Boldt, E. A , , & Swank, J. H. 1982, ApJ, 255, 1 I I
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Astron. Nachr./AN 324, No. S I , 41-46 (2003)/ DO1 10.1002/asna.2003851IS
Magnetic field in the Galactic Centre: Rotation Measure observations of extragalactic sources Subhashis Roy * I , A. Pramesh Rao
I,
and Ravi Subrahmanyan2
' National Centre for Radio Astrophysics (TIFR), Pune University Campus, Post Bag No.3, Ganeshkhind, Pune 41 1 007, India.
' Australia Telescope National Facility, CSIRO, Locked bag 194, Narrabri, NSW 2390, Austrd~ia Abstraet. We have estimated the Faraday Rotation Measure (RM) towards 45 extragalactic sources seen through the central - 6 O < 1 < 6", - 2 O < h < 2" region of the Galaxy using the ATCA and the VLA. This is the first systematic study of RM i n the central kpc of the Galaxy and has lead to the detection of a large scale magnetic field in the region. The magnetic field does not undergo any reversal of sign across the rotation axis of the Galaxy, which is consistent with the bisymmetric spiral model for the magnetic field in our Galaxy. If this component prevails over the central 2 kpc region of the Galaxy, and the average electron density in the region is taken to he 0.4 cm-:', then the line of sight component of this regular magnetic field is estimated to he 0.7 p G . In addition to this large scale field, there are also small scale fluctuations, whose coherence scale length is estimated to he -20 pc. If the scattering material is unilbrmly distributed in the GC. then we estimate the random component of the field to have a strength of 10 pG. On the other hand. if the scattering medium is clumpy, the random component will have similar strength as the regular component. Constraining the filling factor of [he ionised medium is required for estimating the random component of the magnetic field from the R M data.
1 Introduction Magnetic fields can be strong enough to have a significant influence on the dynamics and evolution in the central region of our Galaxy. Magnetic pressure can contribute significantly to the overall pressure balance of the interstellar medium (ISM) and can even influence the distribution of gas (Beck et al. 1996). Strong systematic magnetic fields in the Galactic Centre (GC) region are believed to be responsible for the creation and maintenance of the unique non-thermal filaments (NTFs) (Morris & Serabyn 1996, and references thercin). Measurement of the field geometry and its strength in the central part of the Galaxy is important for estimating its effect on the G C ISM and discrete objects in the GC. There are no estimates of the magnetic ficlds in the inner 5 kpc region of the Galaxy (Davidson 1996) except in the central 200 pc region which are based mainly on observations of non-thermal filaments (Yusef-Zadeh & Morris 1987; Anantharatnaiah et al. 1991; Gray et al. 1995; Yusef-Zadeh et al. 1997). The R M estimated along these NTFs are -1 000 rad m-2 (Yusef-Zadeh & Morris 1987; Gray et al. 1995; Yusef-Zadeh et al. 1997). From the pressure balance argument, (Yusef-Zadeh & Morris 1987) derived a magnetic field strengths of about 1 mG in these NTFs. Previous studies of NTFs have shown that the field lines are oriented along the length of the NTFs. Since, all the NTFs found within a degree from the GC arc oriented almost perpendicular to the Galactic plane, it suggests that the field lines in the surrounding ISM are also perpendicular to the Galactic plane (Morris & Serabyn 1996). However, in the NTF Pelican (G358.85+0.47) located beyond a degree from the GC, the field lines are almost parallel to the Galactic plane. This may indicate that the field lines change their orientation beyond one degree from the G C and become parallel to the disk, as usually seen in the rest of the Galaxy. However, it should be noted that if the * Suhhashis Roy: e-mail:
[email protected] 02001 WILEY-VCH
Verlap GmhH C To K(iaA. W a n h a m
S. Rov et al.: Magnetic field in the Galactic Centre
42
NTFs are manifestations of a favourable local environment (Shore & Larosa 1999) (much higher magnetic field), inferences drawn from these observations can be misleading. While Zeeman splitting of spectral lines directly measure the magnetic field, this method is sensitive to high magnetic fields. Therefore, estimates of mG magnetic field based on Zeeman splitting (Killeen et al. 1992; Yusef-Zadeh ct al. 1999) of the HI or OH lines towards the GC could be atypical, being from local enhancement of field (e.g., near the core of high density molecular clouds) To estimate any systematic magnetic field in the region, it is necessary to use an observational method, which is sensitive to the large scale field. The Faraday Rotation Measure, which is the integrated line of sight (los) magnetic field weighted by the electron density (RM=0.81 x Jn,B,ldl, where n, is the electron density, Bl1 is the 10s component of the magnetic field, and the integration is carried out along the los), is one such method. If a model for the electron density is available, such observations towards the background extragalactic sources are well suited to estimate the 10s magnetic fields prevailing in the ISM near the GC. However, no systematic studies of the Rotation Measures (RM) towards the extragalactic sources seen through the GC region have been carried out in the past. We have carried out a survey of Rh4 towards 65 suspected extragalactic sources seen through the central -6"< 1 < 6", -2"< b 0l -. 199% ApJ, 505,715 Moms,M. &-Serahyn, E. 1996, ARA&A, 34,645 Saikia, D. J. & Salter, C. J. 1988, ARAkA, 26, 93 Shore, S. N. & Larosa, T. N. 1999, ApJ, 521,587 Simard-Normandin, M. & Kronherg, P. P. 1980, ApJ, 242, 74 Taylor, J. H. & Cordes, J. M. 1993, ApJ, 41 1, 674 Yusef-Zadeh, F. & Morris, M. 1987, ApJ, 322,721 Ynsef-Zadeh, E, Roberts, D. A., Goss, W. M., Frail, D. A,, & Green, A. J. 1999, ApJ, 512, 230 Yusef-Zadeh, F., Wardle, M., & Parastaran, P. 1997, ApJL, 475, L119 Zoonematkermani, S., Helfand, D. J., Becker, R. H., White, R. L., & Perley, R. A. 1990, ApJS, 74, 18 I
Astron. Nachr./AN 324. No. S 1.47- 5 1 (2003)/ DO1 IO.1002/asna.20038S024
Study of the Nuclear Bulge region of the Galaxy K. S. Baliyan’ I, S. Ganesh’, U.C. Joshi’, 1.S. Glass’, and T. Nagata’
’ Physical Research Laboratory, Ahmedabad-380 009, India ’ South African Astronomical Observatory, South Africa ’ Nagoya University, Nagoya, Japan
Key words Near Infrared, Milky Way Galaxy, Nuclear Bulge, Photometry, Extinction. PACS 04A25 A near infrared survey of the inner 3OOpc of Nuclear Bulge region of the Milky Way is being carried out as a core program of the SIRIUS Camera mounted on the IRSF telescope at the South African Astronomical Observatory, Sutherland. The SIRIUS camera has three I K x 1K detectors for simultaneous imaging in the J, H and Ks bands. With pixel scale of 0.45” and good seeing most of the time, these observations present the deepest views of a large area of the Nuclear Bulge. The aim of this survey is to overcome the incompleteness and confusion limited nature of the undersampled near infrared surveys, such as DENIS and 2MASS in order to get better estimate of the extinction in these lines of sight and distinguish between
various stellar populations. Preliminary results from the survey demonstrate the capability of the camera and photometric procedures for the crowded fields.
1 Introduction The Galactic Center region of our Galaxy provides the best opportunity to study the astrophysical processes i n galactic nuclei with high spatial resolution. The study of the inner regions of the Galaxy is also very important for understanding the structure, dynamics, kinematics & energetics of the Milky Way, as well as galactic evolution and star formation processes (Serabyn & Morris 1996, Mezgcr, Duschl & Zylka 1996, Ellis 2001, Launhardt et al. 2002). Due to the large extinction, studies of the Galactic structure based on optical observations have been restricted to high Galactic latitudes. However, the longer wavelenglhs observations, including near infrared, can penetrate the relatively high extinction at low Galactic latitudes towards the nuclear bulge. In the recent past there have been several NIR surveys, i.e. DENIS (Epchtein et al. 1997), 2MASS (Skrutskie et al. 1997), etc, which either lack depth or suffer from confusion due to high source densities in the Galactic Center Region. Recently Schultheis et al. (1999) prepared a map of the interstellar extinction towards the inner Galactic Bulge using DENIS data and reported Av> 25 with a clumpy, inhomogeneous nature. However, the J band data i n DENIS is undersampled in this region of high extinction. A large number of Ks sources do not have counterparts in 1 & J in DENIS. The situation has not improved much with the availability oE 2MASS data. To overcome these problems, and to gain a better understanding of the distribution of stellar populations in the inner bulge region, we are carrying out a deep imaging survey of this region i n J, H bz Ks bands with particular emphasis on the fields covered by the ISOGAL survey at 7 p n and 15 p m discussed here by Baliyan et al. (2003) and the X-ray survey by thc Chandra observatory. The deep imaging survey in J, H & Ks bands was carried out using IRSF telescope of the South African Astronomical Observatory at Sutherland during June-July 2002. These observations have been planned so as to reach the tip of the RGB at the distance of the Galactic center. Another key ingredient for studying the spatial structure of * Corresponding author: e-mail:
baliyanQprl.res.in. Phone: +9l 796302 129, Fax: +91 796301 502
02003 WILEY-VCH VrrLtg Grnblj L Co
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the Galactic bulge is deriving accurate distances to the stars in these regions. Such information will help identify foreground and background sources. In this paper here we present fust results from this near infrared survey underlining the capabilities of the SIRIUS camera and the photometric procedures adopted in this work.
2 Thesurvey Our survey of the inner 300 pc of the bulge region within [ I ( = 1.5 deg, (b( = 0.5 deg is carried out using the IRSF (InfraRed Survey Facility) 1.4 m telescope at Sutherland, South Africa. The IRSF is jointly operated by the SAAO, South Africa and Nagoya University, Japan. The telescope is equipped with a three channel camera (Nagashima et al. 1999) known as SIRIUS (Simultaneous InfraRed Imager for Unbiased Survey), capable of imaging in J ( I .25 pm), H ( I .63 p m ) & Ks (2.14 pm) bands, simultaneously with 1024x1024 pixels array detectors. The FOV is 7.8’x7.8’ with a scale of 0.45 arcseclpixel. The region covered in the present survey is shown in Fig. 1 where each square rcpresents the center of a mosaic of 3x3 images. Righl ascension 1:
3
46:OO
44:OO I
I
I
30:OO
-29:oo:oo 2 L
-C .0
0
30:OO
-3O:OO:OO
,
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Fig, 1 Image of the Inner Milky Way showing the surveyed region.
3 Observations We performed imaging observations for the inner Galactic Bulge at near infrared wavelengths during June 25-July 1 & July 9-15, 2002 using the InfraRed Survey Facility (IRSF) at SAAO. The observations were conducted under good seeing (better than 1.3 arcsec) conditions most of the time with airmass values between 1.0 and 1.5. For each field, 10 dithered images were taken (one at the center and 9 on a circle around it) with 0.1 and 5 secs exposures. Towards some directions we have also taken 10 second exposures
Astron. Nachr./AN 324, No. S 1 (2003)
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to reach fainter levels of completeness in the highly extincted J band. Observations were also made for
sky, dark, twilight flat fields and on a set of standard stars (Persson et al. 1998) for data reduction and calibration. Observations for sky were taken towards a dark nebula not far from the object of interest. Some fields were also observed towards low extinction regions symmetrically placed perpendicular to the galactic plane to serve as reference.
4 Data Reduction and Analysis The images were reduced using the IRAF package along with scripts written in-house. Standard procedures of near-infrared array image reduction are adopted, including dark subtraction, sky subtraction and Rat fielding. The SIRIUS Pipeline (Nakajima, private communication) was used for combining the 10 dithered images into a single image to obtain a higher signal to noise ratio for each position. Since the Nuclear Bulge is highly crowded, special care has to be taken while doing photometry. We noticed that using IRAF’s DAOPHOT package for source detection and photometry resulted in a large number of residual sources. We, therefore, performed source detections and photometry using Christophe Alard’s (Alard 2000) software which applies PSF fitting to the image after dividing it into several regions. Thus it uses space varying kernels. This procedure results in far less number of residuals than what is obtained wirh IRAF/Daophot. The limiting magnitudes in J, H and Ks are 17.2, 17 & 16, respectively. For the purpose of this paper, the measured fluxes were calibrated, for preliminary results, using the 2MASS photometry for corresponding obiects.
5 Sample Results & Discussions Here we show some sample results obtained from the 8’ x 8’ colour image (Fig. 2) of the Galactic Center region. The image is coniposcd of J (blue), H (green) and Ks (red) band images. The emission at near infrared wavelengths is much less extincted as compared to optical and is mainly from the photosphere of late spectral type, evolved stars. The central bright star cluster is clearly visible in the image. The image itself indicates the non-uniform, highly varying nature of the extinction. It is easy to see the filamentary and clumpy, inhomogenous distribution of molecular material. The J band shows maximum extinction in the near infrared, decreasing towards longer wavelengths. The sources seen in this band are either blue foreground sources or very bright stars o f the central stellar cluster. The H and Ks band images show overcrowding of the sources in this region. Some morphological details of these molecular clouds is also apparent. The source extraction results i n more than 7500 sources in Ks. The magnitude values used in these preliminary results are calibrated with 2MASS data as mentioned. The Ks versus (J-Ks) colormagnitude (CMD) and (J-H) versus (H-Ks) color-color diagrams are plotted in Figs. 3 and 4, respectively, to distinguish between stellar populations. Only the stars detected in all three bands are used in these plots. The sharp cut off in the Ks versus J-Ks CMD is due to J band limitations. The CMD also shows foreground bright stars (to the Icft) and a small number of highly reddened sources towards left. I t is also possible that there are some massive high luminosity bluer, young stars hidden by the large population of red giants (which are easily detectable in NIR). The color-color diagram shows that stars arc distributed in a much wider range than one finds in less extincted regions. A large number of thc sources right to the main concentration indicates to thc large amount of circumstellar material. The region, therefore, harbours sources belonging to various populations. 11 is to be noted that our survey is deeper by more than 2-magnitudes when compared to 2MASS & DENIS i n the corresponding bands. We are therefore reporting measurements of a large number of sources (fainter than Ks o f 14 mag) for the first time. We also imaged a field lying at higher latitude (I=0.0, b=I.O) with much lower extinction (Av 6, see for example Omont et al. 1999) Tor refercncc purposes. Towards this field, with 50 secs (10 images of 5 secs exposure average combined) integration, we detect more than 18000, 13000 & 12000 sources in J,
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K. S. Baliyan et al.: Near Infrared Survey of the Nuclear Bulge
50
Fig. 2 J, H, Ks colour composite image ~ ( 8 ’ x 8 ’of ) the Galactic Center Region centered at GC
H & Ks bands with more than 9000 detections common to ail three bands. The CMD of this region (Fig 5 ) shows well defined giant branch, red clump and foreground sources in contrast to the central region where a large scatter is observed. Also notice that the CMD of GC rcgion shows stars distributed in a much wider range in colour due to varying and non-uniform extinction across the field compared to that in this reference field. The information on the distribution of various classes of the objects is very useful to get a handle for estimating the extinction towards the Galactic Center and distances to the bulk of the sources in these lines of sight.
6 Conclusions The present work reports on a deep near infrared survey of the inner 300 pc of the nuclear bulge region of the Milky way using SIRIUS camera mounted at the I .4 M lRSF telescope of the SAAO. The preliminary results show unprecedented view of the Galactic Bulge region in the J, H and Ks hands, deeper by more than two magnitudes as compared to DENIS and 2MASS. Therefore, we report the measurement of a large
Astron. Nachr./AN 324, No. SI (2003)
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Z
I
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.
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J-Ka
Fig. 3 Ks versus J-Ks CMD showing sharp cut off on the lower right due to high extinction at J.
51
0
1 H-K,
2
3
Fig. 4 J-H versus H-Ks colorcolor diagram for the Galactic Center Region
Fig. 5 Ks versus J-Ks diagram for the reference field (P = 0, b = 1).
number of sources for the first time. Color-magnitude and color-color diagrams are plotted for sources i n
the sampled region. Acknowledgements This work was funded by the Department of Space, Govt. of India. The observation time at IRSF facility awarded by SAAO is gratefully acknowledged. KSB & SG are thankful to SAAO for hospitality at Cape Town and to Observatory staff for help. We also express our thanks to C. Alard for his codes, M. Schultheis and A. Omont for useful discussions and the 2MASS & DENIS teams for the respective data. ISOGAL provided the background image of the inner galaxy used in the figure 1. Partial funding by DST enabled KSB to attend the GC-2002 conference in Hawaii.
References Alard, C. 2000, A&AS, 144,363 Baliyan, K.S., et al. 2003, these proceedings Ellis, R.S. 2001, PASP 113, 515 Epchtein, N., et al. 1997, Messenger 87, 27 Launhdrdt, R., Zylka, R., Mezger, P.G. 2002, AXrA, 384, 112 Mezger, P.G., Duschl, W.J., Zylka, R. 1996, A&AR, 7, 289 Nagashima, C., et al. 1999, in, Srur Formarion 1999. Ed. T. Nakamoto, Nagama: Nobeyama radio Observatory), 397 Omont, A., et al. 1999, A&A, 348,755 Persson, S.E.,Murphy, P.C., Krzeminski, W., Roth, M. & Rieke, M.J. 1998, AJ, 116, 2475 Serabyn, E. & Morris, M. 1996, Nature, 382, 602 Skrutskie M.F. et a1 1997, in The impact oflurge .scale Near IR sky survry, Eds. Garzon F, Epchtein N B Omont A, K1uwer:Netherlands. 117
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Astron. Nachr./AN 324, No. SI, 53 -57 (2003) / DO1 10.1002/asna.200385025
A morphological Study of the Galactic Inner Bulge Kiran S. Baliyan*', Shashikiran Ganesh', Umesh C. Joshi', Ian S. Glass', Mark R. Morris3, Alain Omont4, Mathias Schultheid, and Guy Simon5 I
' '
Physical Research Laboratory, Navarangpura, Ahmedabad, India South African Astronomical Observatory, Cape Town, South Africa Department of Astronomy, UCLA, USA Institut d' Astrophysique de Paris, Paris, France Observatoire de Paris, Paris, France
Key words Galactic Inner Bulge, mid infrared, stars, molecular clouds PACS 04A2.5
A sizable region of the Inner Galaxy was observed during the ISOGAL survey using ISOCAM narrow band imaging at 7 and 15 pm with 3 arcsec resolution. Due to very low, though non-negligible, extinction at mid-infrared wavelengths, this window is very suitable for exploring the morphological features in the inner Galaxy. The images from this survey display spectacular mid-IR emission and absorption features and highly crowded star fields. The starforming regions of Sgr B1 and C stand-out at mid-IR wavelengths. Also very obvious are quiescent molecular clouds which are seen as dark regions in the mid-infrared images. The diffuse mid-IR emission is from ionixd gas, warm dust (mainly at 15 pm) and PAH (7 pm). The emission from point sources is mainly due to circumstellar dust, but at 7 pm stellar photospheres also contribute to it. As a first step to understand this complex ensemble of sources, we compare these images with previously known information at other wavelengths particularly the VLA 90 cm map. A good correlation between the mid-IR starforming regions and their thermal counterparts in the 90 cm VLA map is seen.
1 Introduction Attempts at understanding the inner Milky Way remain as a forefront research area with several good reviews on the subject of the Galactic Inner Bulge and the nuclear regions are such as by Mezger et al. (1996),Morris & Serabyn (1996) and Launhardt et al. (2002). The MSX survey (Price et al. 1997, Egan et al. 1998) of the Galactic Plane covered the entire plane of the Galaxy. The ISOGAL Survey has revealed a new view of the Inner Milky Way at mid-infrared wavelengths (Omont et al. 2003). This survey at 7 and 15 p m used the ISOCAM camera onboard the Infrared Space Observatory. With a pixel-field of view of 6" (3" in some of the more crowded regions) and sensitivity better than IOmJy, the survey resulted i n a view that is better by two orders of magnitude in spatial resolution and sensitivity as compared to IRAS at similar wavelengths. With its smaller pixel sizes, the ISOGAL images are of a superior resolution compared to MSX. A first look at the results in the disk of the Milky Way was presented by PCrault et al. ( 1996). Subsequently Glass et al. (1999) and Omont et al. (1999) provided a good insight of the relatively outer regions of the Bulge of the Galaxy at mid-infrared wavelengths. They standardized the procedures for the analysis of mid-IR data from the survey (in combination with near-infrared ground based observations) in thc presence of low interstellar extinction (see also Ojha et al. 2003). Many of the stars detected by ISOGAL are late M-type giants on the asymptotic giant branch(AGB) in the Galactic Bulge and central disk but the most numerous class of sources detected at 7 p m are red giants with luminosities just above or close to the RGB tip with weak mass-loss rates. * Corresponding author: c-mail:
[email protected] K. S. Baliyan et al.: A morphological Study of the Galactic Inner Bulge
54
Interstellar extinction poses a major difficulty at optical wavelengths in studies of the inner regions of the Galaxy. The very high and non-homogeneous extinction in these regions (see for example Catchpole et al. 1990, Schultheis et al. 1999) is significant at near infrared wavelengths and non-negligible even at mid-infrared wavelengths along some lines of sight. ISOGAL observations have been used by Hennebelle et al. (2001) to provide constraints on the interstellar extinction curve at mid infrared wavelengths towards the Inner Bulge. Indeed the ISOGAL images present the deepest and clearest views (over a large area) of the Nuclear regions of the Galaxy till date. Version 1.0 of the five wavelength ISOGAL-DENIS point source catalog is now published (Schuller et al. 2003) with discussion of the scientific results (Omont el al. 2003). In this work we present some of the ISOGAL survey images and compare them with already known information at other wavelengths. The images exhibit a good correlation between locations of the dark clouds in mid-IR and continuum emission at millimeter wavelengths. These mid-IR images also confirm the asymmetric distribution of the stellar sources, dust and gas in this region.
2 Observations Observations of the Galactic Nuclear Bulge regions (0.80' > f! > - 1.2", -0.4" < b < 0.25" were made with the 60cm telescope and the camera ISOCAM (Cesarsky et al. 1996), on hoard the Infrared Space Observatory under the ISOGAL program. While most of the ISOGAL observations employed 6" pixels, in the dense regions of the Nuclear Bulge 3" pixels were used. Narrow band filters, LW5 at 7 p m and LW9 at 15 pm, were used for the observations discussed here. Further details of the observations and the data reduction techniques are discussed elsewhere (Schuller et al. 2003, Ganesh et al. in prep.)
3 Discussion of the images
'grB2
I
Sgr B l
G0.25+0.02
GC
Fig. 1 A mosaic of the 7 pm ISOGAL images superimposed VLA 90 cm contours. This figure appears in colour in the electronic version
Figure 1 shows the mid infrared view of the Inner Milky Way constructed from the ISOGAL observations at 7 pm. Very evident are spectacular emission features including Sgr B and Sgr C. Superimposed on the emission features are also regions marked by strong absorption at 7 pm. One well studied (see for example Lis & Carlstrom, 1994) giant molecular cloud G0.25+0.02 is marked in Figure 1 where it
Astron. Nachr./AN 324. No. S1 (2003)
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appears as a prominent dark absorption feature. Apart from diffuse emission, we also see largc number of stars resolved for the first time by the superior dctcctors+telescope combination of the [SO. Most of these were unresolved by IRAS and some also by MSX. We note that with relatively large pixel size, some of the bright ISOGAL sources could also be unresolved compact clusters. For example, the Arches Cluster would be covered by a single 6” pixel of ISOCAM. However, many of such young objects appear as slightly extended sources (Schuller 2002). It is clear that the number density of stars on both sides of the Galactic Centre is non-uniform. Towards the negative longitudes we have larger number of stars as compared to that at positive longitudes. This asymmetry is perhaps due to the higher extinction towards the positive longitudes where there are relatively larger number of dark globules and patches indicating giant molecular clouds. Superimposed on 7 p i image in Figure 1 are contours from the VLA obscrvations at 90 cm (LaRosa et al. 2000). The correspondence of the strong emission at 90 cm with extended 7 pm emission is seen very clearly lor the Sgr B 1 starforming region and also the Sgr C region. Towards Sgr B2 however, the 7 pm mosaic shows a lot of absorption by intervening molecular cloud layers as discussed above.
Fig. 2 An enlarged view of the Sgr B region at 7 pi from ISOGAL superimposed with VLA 90 cm contours. Orientation is in Galactic Coordinate system for all ligures.
Since the VLA is sensitive to both the thermal and non-thermal emission, interesting correspondences are seen with mid-IR features. Several ridge-like extended emission features appear correlated with VLA
56
K. S. Baliyan et al.: A morphological Study of the Galactic Inner Bulge
flux contours. Above the Sgr Bl & B2 complex, towards positive latitudes, a dark cloud, with low 90 cm flux, appears to he relatively quiescent with little, if any, star formation activity.
Fig. 3 Magnified view of the Sgr C region from figure 1
Among the extended emission features, the star forming regions, Sgr B1 (see Figure 2) and Sgr C (Figure 3), stand-out at mid-IR wavelengths and exhibit interesting correspondences with structures at 90 cm seen in the images obtained by the VLA. In the interest of keeping the figures clear and to display the large number of mid-infrared point sources in proper contrast, we do not overlay astrometric grids on these figures. In the case of Sgr B 1, a computational artifact in the VLA contours runs across the molecular cloud but the correspondence between the contours and the 7 pm extended emission features is clearly
Astron. Nachr./AN 324, No. S 1 (2003)
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seen. High extinction, even in the mid-infrared, apparently hides from view the strong VLA source Sgr B2. There are a fcw mid-infrared poinl sourccs scattered around the actual peak of radio emission. In Figure 3, some 7 p m sources appear aligned along the non-thermal filament(NTF) associated with the Sgr C. Whether there are any correspondences between mid-IR sources and radio NTFs is however, still an open question, as the apparent alignment could be purely coincidental. Further discussion of these images including comparison with molecular data will be published elsewhere (Ganesh et al. 2003). An electronic-version-only figure' shows a color image constructed by mapping mosaic of 7 pin to green and 15 pm to red. The result shows colors of the point sources and the extended emission. At 15 pm, there is a ridge-like feature extending Crom (I,b =0.05, -0.1) to (l,b=0.22,-1.0), touching b=-0.18 midway, which is almost absent in the 7 pni image. It appears like a compressedlshocked shell perhaps driven by nearby supernova remnant. This feature seems to be in the same line of sight as the base of the radio arched filament. MSX images show this feature as a super bubble, perhaps powered by the Quintuplet cluster. Acknowledgements This work is supported by the Indo-French Centre for the Promotion of Advanced Research under project 1910-1. Financial support from Department of Space & Department of Science & Tech, Govt. of India made it possible to present this work at the GC-2002 conference at Hawaii. The VLA 90 cm data were obtained from the web pages of Lazio (http://rsd-www.nrl.navy.mil/72 I3/lazio/GCatlaa/)
References Catchpote, R. M., Whitelock, P.A., Glass, I.S. 1990, MNRAS, 247,479 Cesarsky, C.J., et al. 1996, A&A, 315, L32 Egan, M.P., Shipman, R.F., Price, S.D., Carey, S.J., Clark, F.O., Cohen, M. 1998, ApJ, 494, 199 Ganesh, S., et al. 2003, A&A, in preparation Glass, I.S., et al. 1999, MNRAS, 308, I27 LaRosa, T.N., Kassim, N.E., Lazio, J.W., Hyman, S.D. 2000, AJ, 119, 207 Launhardt, R., Zylka, R., Mezger, P.G. 2002. A&A, 384, 1 12 Lis, D. C. & Carlstrom, J.E. 1994, ApJ, 424, 189 Mezger, P.C., Duschl, W.J., Zylka, R. 1996, A&AR, 7,289 Moms, M.R., & Serabyn, E. 1996, ARA&A, 34,645 Ojha, D.K., et al. 2003, A&A, 403, 141 Omont, A,, et al. 1999, A&A, 348,755 Omont, A,, et al. 2003, A&A, 403, 975 PCrault M., et al. 1996, A&A 315, L165 Price, S.D., et al. 1997, IAU Symp. 179, 1 IS Schuller, F. 2002, PhD Thesis (University of Paris). Schuller, F., et al. 2003, A&A, 403,955 Schulthcis, M., et al. 1999, A&A, 349, L69
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Astron. Nachr./AN 324, No. SI, 59-63 (2003) / DO1 10.1002/asna.200385026
Warm molecular gas, dust and ionized gas in the 500 central pc of the Galaxy N.J. Rodriguez-Fernandez * I , J. Martin-Pintado2, A. Fuente3, and T.L. Wilson4 ' LERMA, Ohservatoire de Paris,61, Av de I'Observatoire, 75014 Paris, France
' Instituto de Estmcturd de la Materia, CSIC, Serrano 121, 28006, Madrid, Spain
' Observatorio Astron6mico Nacional. LGN, Apdo. 1143, 28800 Alcali de Henares, Spain Max-Planck-Institutfur Radioastronomie,Auf dem Huge1 69,53121 Bonn, Germany
Key words ISM: lines - ISM: Infrarcd - Galaxy: center
PACS 04A25
We present infrared and millimeter ohservations of molecular gas, dust and ionized gas towards a samplc of clouds distributed along the 500 central pc of the Galaxy. The clouds were selected to investigate the physical state, in particular the high gas temperatures, of the Galactic center region clouds located far from far-infrared of thermal radio continuum sources. We have found that there is ionized gas associated with the molecular gas. The ionizing radiation is hard 35; 000 K) but diluted due to the inhomogeneity of the medium. We estimate that- 30% of the warm molecular gas observed in the Galactic center region clouds is heated by ultra-violet radiation in photo-dissociationregions.
+
1 Introduction The interstellar medium in the central SO0 pc of the Galaxy (hereafter Galactic Center, GC) is mainly molecular gas. The molecular clouds in the GC exhibit an extended gas component with high temperature ( IS0 K). On the contrary, the dust temperature is lower than 30 K. The large line widths of the molecular lines, the high gas phase abundance of molecules linked to the dust chemistry, and the difference between gas and dust temperature suggest that some kind of shocks could he responsible for thc high gas temperalures of the molecular gas (Wilson et al. 1982, Martin-Pintado et al. 2001). The possible influence of radiation in the heating of the molecular gas is usually ruled out due to the lack of far infrared and thcrma1 radio continuum sources in the G C others than the well known H I I regions associated with the Sgr complexes (A-E) or ionized nebulae like the Sickle To investigate the heating of the molecular clouds in the GC we have studied a sample of 18 clouds located all along this region. The clouds were selected as molecular peaks located far from far infrared or radio continuum sources. Thosc sourccs were observed with the spectrornetcrs on board the lnfrural Spncc~ Observatory (ISO) and with the IRAM 30-in telescope. The data obtained with IS0 are ohservations of the lowest H2 pure-rotational lines, dust continuum spectra from 40 to 190 pm and a number of fine-structure lines from neutral atoms or ions (e.g. 0 I 6.1 pm, C II IS8 p m , Ne II 12 p m , 0 III 52 pm). With the IRAM 30-m antenna we have observed C"0, " C 0 , H 3 S a and H41 a. I n this paper we will review the results already published by Rodriguez-Fernandea et al. (2001a, 2001b. hereafter RFOla, RFOIb, respectively) and present some new results that will b e extended elsewhcrc (Rodriguez-Fcrnandczet al. 2003, RF03).
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* e-mail: nemesio.rodriguezQobspm.fr, Phone: +33 140512061, Fax: +33 140512002 @ 2003 WILEY-VCH Vsiiag tinihH B Co. KGdA, Wcinliciin
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N. J. Rodriguez-Fernhdezet al.: Warm molecular gas. dust and ionized gas in the 500 central uc of the Galaxv
2 Warm molecular gas Before ISO, the warm gas component of the GC molecular clouds was mainly studied by means of NH3 observations. The multilevel study of Huttemeister et al. (1993) showed that the temperature structure of the GC clouds can be characterized with two gas components at different temperatures: a cold gas component with a temperature close to that of the dust ( N 20 K) and a warm gas component whose temperature ranges from 100 to 250 K. However, since the abundance of ammonia is known to vary significantly is difficult to estimate the column density of warm gas in the GC clouds. We have, for the first time, obtained the column density of warm gas in the GC clouds by observing Hz pure rotational lines (RFOla). Columns 3 and 4 of Table 1 show the total column density and the temperature of the warm gas, respectively. We have also estimated the HZ density and the total column density of molecular gas in these clouds by observing 13C0 and C"0 and doing aradiative transfer analysis for kinetic temperatures between 15 and 200 K. The results are shown in columns 1 and 2 of Table 1. The fraction of warm H? to the total H2 column density as traced by CO varies from source to source but is 30% on average. As discussed in RF0 1a, it is difficult to explain the large column density of warm gas in the GC clouds. Several low velocity (- 10 Km s-l) C-shocks, photo-dissociation regions (PDRs) or both should be present in the line of sight. Comparing the energy of the turbulent motions in the GC clouds with the cooling by H2 (which at the moderate density of the GC clouds is comparable to that by CO), one finds that dissipation of supersonic turbulence could account for the heating of the warm Hz. However, with the available data is not possible to rule out heating in PDRs, indeed the observed temperature gradient in the GC clouds can be appropriately reproduced in a context of a PDR (RFOI a).
3 Dust temperatures The dust continuum emission peaks at wavelengths of 100 to 80 pm for all but two sources, whose spectra peak at SO-60 pm. The dust luminosity in the observed range is listed in column 5 of Table 1. It is not possible to fit the spectra with just one grey body. Thus we have used a model with two grey bodies like that described in section 2 of Goicoechea et al. (2003). Figure 1 shows the data and the grey bodies fits for two sources. To explain the emission at long wavelengths it is needed a dust component with a temperature of 15 to 18 K. The temperature of the warmer component varies from source to source from 26 to 39 K. Due to the uncertainties in the dust emissivity it is not easy to determine a total column density of dust (on the contrary, temperatures are almost independent on the dust emissivity). Nevertheless, the column density of dust with temperatures higher than SO K is less than 500 times lower than the column of dust at 15-35 K.
50
1 00 150 lung onda (micros)
200
Fig. 1 Left panel: Dust continuum emission towards M+0.21-0.12 (black solid triangles) and grey-body fit (green) with two temperature components ar 39 (red) and 16 K (blue).Right panel: same for M+0.76-0.0.5 with a 24 K (red) and 1.5 K (blue) components.
61
Astron. Nachr./AN 324, No. S I (2003)
Table 1 Physical parameters derived from the observations. See texi [or explanation. Number in parentheses are errors of the last significant digit. Typical errors of TI and Tz are 5 and 1 K, respectively.
M-0.96+0.13
3.5-4
0.6-1.1
I.lO(9)
157(6)
2.9
30
15
547.8
52.5
...
M-0.55-0.05
3.8-4.4
4.3-6.0
2.7(3)
135(5)
12.3
31
16
547.5
57.1
...
M-0.50-0.03
3.4-3.7
2.4-3.0
2.3(2l
135(4)
10.5
34
17
548.0
52.9
...
M-0.42+0.01
4-4.5
2.1-3.4
1.03(8)
167(6)
9.8
30
16
548.0
530
...
M-0.32-0.19
>3
1.1-2.2
1.03(51
188(5)
7.7
35
16
547.6
4.4
34
M-0.15-0.07
3.7-4.1
6.6-8.4
2.6(4)
136(6)
...
._.
._.
547.6
...
...
M4.16-0.10
3.8-4.2
3.7-4.9
1 17(13l
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178(5)
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586
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4 Radio recombination lines To study the possible presence of ionized gas associated with the GC molecular clouds we have observed the H41a and H35a recombination lines using the IRAM 30m antenna (RF03). However, we have no1 detected any of the lines in any of the sources. From our 3 a limits to the line fluxes we have determined upper limits to the flux of Lyman continuum photons in the 30m beam (column 8 of Table I ) . The comparison of thosc limits with stellar atmospheres models rules out the presence of stars of an earlier type than B0.5 and effective temperatures higher than 32,000 K. N
5
Fine structure lines
We have detected lines of atoms and low excitational potential (< 13.6 eV) ions like 0 I 63 pm, C II I58 p m or Si I I 34 p m in all of the sources. In most of them we have also detected N 11 122 pm or Ne 11 12 p m (excitational potential of 14 and 21 eV respectively). In 1 1 of the 18 sources we have even detected the 0 I I I 88 p m line (the excitational potential of 0 I I I is 35 eV). Column 9 of Table 1 shows the electron densities of the ionized gas component as derived from the 0 III 52 to 0 III 88 p m lines ratios. The densities, around 10-100 crnp3, are lower than those found in the Radio Arc region (RFOlb) and in Sgr B2 (Goicoechea et al. 2003). Column 10 of Table 1 shows the effective temperatures derived from the N I I I 57 p m to N 1 1 122 p m lines ratios for the threc sources were the N I I I line has been detected. These temperatures have been calculated following the H 11 regions models by Rubin et al (1994). However, as pointed out by Shields and Ferland ( 1 993), the observed lines ratios can even be reproduced with higher
62
N. J. Rodriguez-Femandez et al.: Warm molecular gas, dust and ionized gas in the 500 central pc of the Galaxy
effective temperatures of the ionizing radiation if the ionization parameter is low, i.e., if the medium is clumpy and inhomogeneous and the ionizing sources are located far from the ionized nebulae. These seems to be the case in the Radio Arc region, were we have shown the presence of an extended component of gas ionized by the combined effect of the Quintuplet and the Arches clusters (RFOlb). The radiation can reach large distances due to the inhomogeneity of the medium (in part due to the presence of a large bubble clearly seen in infrared images). Photoionization model calculations showed that the lines ratios observed in this region can be explained with a constant effective temperature of the ionizing radiation but a different ionization parameter for each cloud, consistent with the different distances of the clouds to the ionizing sources. The analysis is more difficult for the clouds located far from the Radio Arc region since the geometry of the medium and the possible ionizing sources are unknown. However, photoionization simulations for 35,000 K are possible if the ionization many lines ratios demonstrate that effective temperatures of RF03). parameter is low (Some of the fine structure lines have been observed in the Fabry-Perot mode. The spectral resolution of this mode (- 30 km s-l) give us the possibility of studying the line profiles of the broad lines from the GC clouds. Figure 2 shows a sample of the lines observed in this mode towards two sources. Taking into account the moderate spectral resolution, the line profiles and the line centers of highly excited ions (like 0 III ) are in good agreement with those of the neutral or low excitational potential ions (like 0 1 and C II ). Furthermore, Fig. 3 shows that the agreement of the line profiles of the weakly ionized gas and the molecular gas is rather good. This fact suggest that the three components are associated, and that radiation could play a role not only in the ionization of the ionized gas but also in the heating of the neutral gas. We have compared the observed far infrared continuum and the C 11 , 0 I and Si 11 lines fluxes with the predictions from PDR models. The power radiated by lines is 0.5% of that radiated by the continuum. While the C II /O I ratio is 5. Plotting both quantities in a plot like that of Fig. 2d of Goicoechea et al. (2003), one finds that the GC clouds exhibit similar properties to those of Sgr B2, with a far ultra-violet field lo3 times larger than the local interstellar radiation field and a hydrogen density of lo3 cmP3. W cmP2)are also consistent The absolute fluxes of the C 11 (- lo-’’ W cmP2)and the 0 1 line (- 2 with the Hollenbach et al. (199 I ) PDRs models predictions.
-
-
-
-
N
6 Discussion: heating and ionization The IRAM-30m “radio view” of the GC interstellar medium seems to confirm the “classical” idea of a mainly neutral and dense gas. However, the global picture arising from the IS0 observations is a complex interstellar medium with associated molecular, atomic, and ionized gas components with decreasing den~ n f ~ lo3, n, 10’-’ cmP3). The ionizing radiation is hard (effective temperatures sities ( n ~ lo4, close to 35,000 K) but diluted due to the large distances from the ionizing sources to the nebulae (- 50 pc in the Radio Arc region). The large distance effect of the ionizing radiation can only be explained if the interstellar medium is very inhomogeneous. This scenario is also necessary to explain the low number of Lyman continuum photons derived from the radio recombination line observations. The radiation must also influence the atomic and molecular phases. As we have seen, the 0 I and C I1 lines are well explained by a PDR with a H density of lo3 cmP3 and a far ultra-violet incident field lo3 times higher than that in the local interstellar medium. The absolute 0 I and C I I lines fluxes predicted by the PDR models are also similar to the observed fluxes arising from the GC clouds. Those models also predict a warm H2 layer with temperatures of 150 K, as we have derived from the pure-rotational lines. However, the total column density of warm Hz predicted by the models is 3 x lo2’ cm-’ while the total lo2’ cm-’. Thus, column density of warm H2 that we have derived from the pure-rotational lines is we conclude that approximately 30% of the warm molecular gas in the GC clouds arises in PDRs in the external layers of the clouds. It is important to note that the discrepancy of dust and gas temperatures only rules out gas heating by gas collisions with hot dust, but it does not rule out ull radiative heating mechanisms. In the external layers 150 K by photo-electric effect on the dust of the proposed PDRs the gas is heated to temperatures of
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Astron. Nachr./AN 324. No. SI (2003) M+O 58-0 13
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-200
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9
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Fig. 2 Fdbry-Perot spectra
-!m
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. . . . . . 7m
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Fig. 3 A comparison of the line profile3 of ‘3CO(l-O) (green) and C II 158 ,urn (blue) towards M+0.24+0.02 (up) and M+0.580.13 (down). Intensity in arbitrary units.
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grains without heating the dust to temperatures higher than 35 K (Hollenbach et al. 1991). In fact, the 30 K dust component in the GC can be associated with the 150 K gas. At least a fraction of the other 7 0 % of warm gas should be heated by shocks. The main evidence for shocks in the G C are the high degree of turbulence revealed by the large line-widths and the high abundance in g a s phase of molecules linked to the dusr chemistry as S O , NH3 or C ~ H S O H(Martin-Pintado et al. 2001), which are easily photo-dissociated i n the presence of ultra-violet radiation.
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Acknowledgements NJK-F has has been supported by a Marie Curie Fellowship of the European Community program “Improving Human Research Potential and the Socio-economic Knowledge base” under contract number HPMFCT-2002-01677. NJK-F acknowledges useful discussions with J.K. Goicoechea.
References Goicoechea, Rodriguez-FernBndez and Cernicharo 2003, these proceedings Hollenhach, D. J., Takahashi, T., & Tielens, A. G. G. M. 1991, ApJ, 377, 192 Huettemeister, S., Wilson, T. L., Bania, T. M . , & Mai-tin-Pintado, J. 1993,A&A, 280, 255 Martin-Pintado, J., Rizzo, J. K., de Vicente, P., Rodriguez-FemBndez, N. J., & Fuente, A. 2001, ApJ, 548, L65 Rodrfguez-Fem2ndez. N. J., Martin-Pintado, J.. Fuente, A., de Vicente, P., Wilson, T. L.. & Hiittemeister, S. 2001n. A&A, 365, 174 Rodriguez-Fernindez, N. J., Martin-Pintado, J., & de Vicente, P. 2001, A&A, 377,63 I Rodriguez-FernBndez 2003 (in prep). Kubin, K. H., Simpson, J. P., Lord, S. D., Colgan, S. W. J., Erickson, E. F., & Haas, M. K. 1994, ApJ, 420, 772 Shields, J. C. & Ferland, G. J. 1994, ApJ, 430, 236 Wilson, T. L., Ruf, K., Walmsley, C. M., Martin, R. N., Batrla, W., & Pauls, T. A. 1982, A&A, 1 15, 185
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Astron. Nachr./AN 324, No. S I , 65 - 71 (2003 ) / DO1 IO.l002/asna.200385063
Prospects for LOFAR Observations of the Galactic Center N. E. Kassim*',T. J. W. Lazio', M. Nord'.2,S. D. Hyman', C. L. Brogan**4,T. N. LaRosa', and N. Duric'
' Code 7213, Naval Research Laboratory, Washington, DC, 20375-5320, USA Department of Physics and Astronomy, University of New Mexico, Albuquerque, NM, 87131, USA
' Department of Physics, Sweet Briar College, Sweet Briar, VA 24595, USA
' National Radio Astronomy Observatory, Socorro, New Mexico, 87801, USA ' Department of Biological and Physical Sciences,Kennesaw State University, Kennesaw, GA, 30144, USA Key words LOFAR, low radio frequencies, Galactic center, nonthermal radio emission, transients, cosmic rays Abstract. Continued improvements in existing low frequency radio interferometers are expected, but limits of sensitivity, angular resolution, and frcqucncy range impose fundamental restrictions which cannot easily be overcome. This has inspired the development of the Low Frequency ARray (LOFAR) which will provide significantly increased imaging power over present low frequency systems. In light of advantages offered by recent low frequency observations of the Galactic center, we consider how LOFAR might impact the field by the end of this decade.
1 Introduction The breakthrough to long wavelength, high resolution imaging (Kassim et al. 1993) has ushered in a quiet renaissance in low frequency radio astronomy making contributions to several areas of Galactic center (GC) research. The 330 MHz VLA image presented in Tucson (Kassim et al. 1999; LaRosa et al. 2000) provided a striking and practical atlas of the inner few hundred parsecs. It also revealed new nonthermal sources including the first parallel nonthermal filament (NTF) (Lang et al. 1999), and the first NTF to exhibit a uniformly decreasing spectral index with length (LaRosa, Lazio, & Kassim 2001). This meeting has seen results from continued improvementsin low frequency imaging. Nord et al. (2003, hereafter N03) presented a higher resolution 330 MHz image nearly tripling the number of NTF candidates and revealing hundreds of small diameter sources. Additional randomly-oriented NTF candidates challenge the model of a simple, ordered GC magnetic field (LaRosa et al. 2003), and radio transients (Hyman et al. 2002) have motivated a focused GC monitoring program (Hyman et al. 2003). The first 74 MHz VLA image (Brogan et al. 2003) provides insights into previously poorly understood sources and affords opportunity to determine the radial locations of GC sources. Continued low frequency improvement at the VLA and GMRT are forthcoming, hut fundamental limitations are imposed by spot frequency coveragc, limited collecting area, and array size. This has inspired the development of the Low Frequency ARray (LOFAR), which will provide significantly increased imaging power. We consider how LOFAR will impact GC research as it becomes operational over the course of this decade. * Corresponding author: e-mail: Narnil:Kassim~~irl.riavy.~i~il, Phone: +01 202 767 0668, Fax: +01 202404 8894 * * NRAO Jansky Postdoctoral Fellow.
(?J 2003 WILEY-VCH V e r l q GinbH & Cu. KGilA, Wwnhnm
N. E. Kassim et al.: LOFAR and the Galactic Center
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2 LOFAR Capabilities LOFAR will revolutionize low frequency radio astronomy by advances in three key areas: 1 ) broad-band frequency coverage from 10-240 MHz; 2) improved angular resolution, by extending maximum baselines to 400 km; and 3) significantly enhance sensitivity, afforded by lOOx increased collecting area over previous or current systems. Towards the GC, nominal LOFAR capabilities will be modified relative to extragalactic observations. Sensitivity will be affected by increased sky noise, but will still significantly exceed current capabilities. Meaningful comparisons for the GC are between the VLA’s sensitivity (- 6 hr) at [330, 741 MHz of approximately [3, 1001 mJy, and LOFAR’s projected sensitivity at [240,74] MHz of [0.1, 31 mJy, respectively. Arcsecond angular resolution will be tempered by frequency and position dependent scattering, and absorption will render certain regions opaque. However, these latter effects can be employed to map the distribution of ionized gas and to determine the radial locations of sources. Furthermore, LOFAR’s broad-band frequency coverage will provide measurements at lower and intervening frequencies currently unavailable. Frequency flexibility is important towards the GC, where it can disentangle competing effects, both intrinsic and extrinsic, which will modify the observed nature of emission and absorption processes. Finally, its -5000 baselines will exceed the baseline coverage at the VLA and GMRT by 1 OX, resulting in much better image fidelity.
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3 Sgr A Complex While scattering and absorption will be significant towards the Sgr A complex at frequencies below 100 MHz, LOFAR can avoid much of these efrects by observing at higher frequencies. Even below 100 MHz, the absorption is patchy, and many regions around Sgr A will be open to investigation. For example, while Sgr A East is visible at 123 MHz but absorbed at 1 10 MHz, portions of the Radio Arc are visible down to 57.5 MHz (LaRosa & Kassim 1985; Kassim et al. 1986; Brogan et al. 2003) and other GC nonthermal sources can be traced to lower frequencies. N
3.1
Sgr A*
The low frequency flux density of Sgr A* anchors continuum spectra predicted by models of the black hole accretion disk and its surroundings (e.g., Yuan et al. 2003). Alternatively, if the spectrum turns over from either extrinsic (e.g.. thermal absorption, Pedlar et al. 1989) or intrinsic (e.g., synchrotron-self absorption, Beckert et al. 1996) propagation effects, this can be used as a further diagnostic of physical processes in the Sgr A environment. NO3 summarize low frequency measurements of Sgr A*, including their possible new 330 MHz detection. If that result is robust, it represents the lowest frequency at which Sgr A* has been detected. However, NO3 acknowledge that the complex effects of confusion require independent verification of their measurement. Scaling the scattering diameter of Sgr A* to 240 MHz, we expect a mean angular diameter of 17”, corresponding to baselines of approximately 15 km. Approximately 50% of LOFAR’s collecting area will be on baselines shorter than 15 km, with an 8-hr (at 4 MHz bandwidth) GC thermal noise limit of 0.2 mJy. Allowing for a further 1OX reduced sensitivity due to possible foreground confusion from intervening regions of the Sgr A complex, we estimate LOFAR could detect or set a 10 d y upper limit on the flux density of Sgr A* at 240 MHz. This is a conservative estimate since LOFAR’s frequency and resolution flexibility, and excellent baseline coverage will significantly mitigate confusion. LOFAR could thus confirm the new VLA result, albeit at a somewhat lower frequency where intrinsic or extrinsic absorption could play a more important role. However, even the existing spectrum of Sgr A* (N03, Figure 7) indicates that any such absorption turnover would have to be very sharp in frequency to hide a detection by LOFAR.
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Sgr A East and West
Pedlar et al. ( 1 989) placed Sgr A West in front of Sgr A East because the thermal gas in the former absorbed the nonthermal emission from the latter. Similar measurements also can be used to constrain H 11 region electron temperatures independent of LTE assumptions (Subrahmanyan & Goss 1996). Measurements at additional frequencies ( w 150 MHz) would constrain the Sgr A West electron temperature as a function of depth in the source, and, by extending measurements to lower frequencies, track the transition from optically thick thermal emission to thermal absorption. Such measurements, also applicable to other thermal GC sources, can provide improved estimates of H 1 1 region filling factors, emission measures, and distances (Kassim et al. 1989). Radio spectral index is a valuable tracer of the relativistic energy spectrum. A high fidelity image of Sgr A East at 240 MHz could be combined with existing higher frequency images for a robust spectral index study. In classic shell-type supernova remnants (SNRs) the relativistic electrons are generated via diffusive shock acceleration and the spectral signatures are well known. While recent X-ray results have reduced the initial controversy over whether Sgr A East is a SNR, there remains some debate on its nature. A spatially resolved spectral index map, made accurate by a broad frequency baseline anchored at low frequencies, will provide a good confirmation of the presumed SNR identity.
4 Nonthermal Sources 4.1
Nonthermal Filaments (NTFs)
NTFs remain mysterious structures, unique to the GC. An improved understanding of them should reveal important insights into the GC magnetic field. The first half dozen NTFs were discovered at centimeter wavelengths and, in retrospect, were perhaps only the brightest and most centrally located NTFs. NO3 have discovered a population of NTF candidates, which if confirmed, would roughly triple the population, suggesting we are probing the tip of a luminosity function. NO3 (their Figure 8) find a luminosity function for the number of NTFs of N x I - " / 5 , where I is the NTF radio brightness at 330 MHz. Moreover, the new candidates were revealed in an image optimized to detect sub-arcminute scale structure. We consider NTF detection at an intermediate LOFAR frequency of 120 MHz, which avoids most free-free absorption, while taking advantage of nonthermal NTF spectra. Assuming 50% of LOFAR's collecting area available to detect structure at this scale, attenuating the sensitivity by -1OX due to sky noise, we estimate I 1.5 mJy beam-' at 120 MHz. Assuming an NTF spectral index of -0.7 implies LOFAR should more than triple again the number of known NTFs, possibly detecting 100 objects at frequencies 2 120 MHz. Given the baseline and frequency coverage which will probe a large range ot NTF parameter spacc, we cxpcct the known source population to be increased significantly. If NTF orientation traces GC magnetic lields, randomly oriented, though generally less bright and shorter NTFs might suggest superposition of a large scale, ordered field upon a smaller scale, tangled component, or it may be related to NTF age. Alternatively, NTFs have been explained as manifestations of local field conditions (Shore & LaRosa 1999). Regardless of the correct explanation, increasing the known population by lifting selection effects influencing current interpretations will clearly benefit understanding of the phenomena.
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4.2
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SNRs
SNR catalogs are affected by severe selection effects and are incomplete at surface brightness levels 8x W m-' Hz-' sr-' (Green 199 I ), which LOFAR will cxcecd hy two orders of magnitude. Gray (1994a) has shown that the number of SNRs in the GC is enhanced, tracing an increased star formation rate. Nonetheless, the census of SNRs ai the GC remains incomplete, both for older SNRs and bright. compact, young SNRs; a more complete census will improve constraints on GC SNlSNR birthrates, statistics, energy input to the interstellar medium (ISM), and enable comparison with progenitor populations.
N. E. Kassim et al.: LOFAR and the Galactic Center
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The usefulness of low frequency observations for revealing new GC SNRs is well documented (Kassim & Frail 1996; Gray 1994b; Roy & Rao 2002; Bhatnagar 2002), and LOFAR's much greater sensitivity and baseline coverage will easily uncover many new SNRs. 4.3
Small Diameter Sources
NO3 also have increased the census of low frequency, small diameter sources by -3X (to 240 sources) compared to LaRosa et al. (2000). While most are extragalactic, a Galactic population has also been identified, including a steep spectrum population of pulsar (PSR) candidates. If confirmed, this would be consistent with the increased star formation rate reflected in the SNR density (Gray 1994a) and the suggestion that the GC may harbor a large PSR population (Cordes & Lazio 1997). Estimates of the luminosity function from NO3 suggest that LOFAR will increase the census to over 1000 sources. With its frequency flexibility, it will be possible to use these and other background sources as probes of ionized gas, through separation of frequency dependent scattering and absorption effects.
5
Transients
Enhanced stellar densities make the G-C a promising region to search for transient radio emission. The strongest evidence is from X-rays, where successively more sensitive observations have pointed towards a high population of white dwarf, neutron star, and black hole binaries with a density peaked towards the GC (Skinner 1993; Fender & Kuulkers 200 I). This is evident in new results from Chandra and XMM-Newton (Wang et al. 2003; Predehl et al. 2003), which reveal a large population of small diameter sources, including >2000 hard X-ray sources likely associated with accreting white dwarfs and neutron stars (Muno 2003). These results imply a large number of X-ray binary systems, known sources of transient radio emission. It is therefore natural to expect a corresponding concentration of radio transients. Unfortunately, given the poor efficiency of conventional radio observations, only a handful of radio transients have been identified. The limitations of previous radio searches have motivated the GC monitoring program reported by Hyman et al. (2002, 2003). Their observations exploit the large 330 MHz VLA field of view (- 2.5") and have revealed two bright transient sources and a few variable candidates. Their work implies radio transients above 50 mJy are either very infrequent ( w one every few years) or have time scales much shorter than a month. However these results are tightly constrained by the sensitivity of the VLA and the impracticality of much more frequent monitoring. Both of these severe selection effects are lifted by LOFAR. LOFAR will be a powerful instrument for studying the bursting and transient universe. The multibeaming design incorporates a dedicated All Sky Monitor (ASM) which could continuously observe the GC for at least several hours a day, efficiently monitoring for transients over a wide spectrum of time-scales. Its hardware buffer will provide a "look back" capability that could be triggered by transient detections at other wavelengths or by LOFAR itself. Therefore LOFAR will enable the first sensitive and unbiased survey for GC transients. While scattering and absorption will limit searches at lower frequencies, searches above 120 MHz. where the sensitivity is enhanced by -1OX over the 330 MHz VLA, should be profitable. In regions more than N 1" from Sgr A*, there will be significant transparency at lower frequencies, enhancing the possibility of discovering a bursting population of coherent, intrinsically steeper spectrum sources. It i s difficult to estimate the number of GC transients LOFAR may detect, but a lower limit comes from considering the number of low mass X-ray binaries (LMXBs). Of the small diameter X-ray sources identified by Sidoli et al. (2001). -25% are LMXBs. If a similar percentage of the mainly unidentified -2000 new Chandra sources are LMXBs, this suggests a population of at least several hundred radio transients. Therefore LOFAR's ability to conduct dedicated, efficient GC monitoring may reveal a potentially large GC radio transient population.
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6 Lower Frequencies Low frequencies (< 100 MHz) offer a dramatically different view of the inner Galaxy, as shown at this meeting by comparison of 74 and 4800 MHz images (Brogan 2003). The difference is due to H 11 regions becoming opaque at low frequencies, which can be used as an important tool for Galactic astrophysics (Kassim 1990). 6.1
3D View of the Galactic Center
The appearance of sources in absorption on low frequency GC maps constrains their properties and radial position. Examples include Sgr A East and West (53), the H II regions Sgr D and Sgr E, placed in front, and behind the GC, respectively, and the appearance of the GC Omega Lobe in absorption, delineating a thermal component (Brogan 2003). These latter results emerge from a VLA 74 MHz image with -20X higher angular resolution and higher sensitivity, respectively, then previously available. LOFAR will push resolution and sensitivity at least 1-2 orders of magnitude further, providing measurements towards many inore sources. Moreover its broad-band antennas will permit observations at successively lower frequencies to smoothly trace the onset of frequency dependent propagation effects. This can be used to investigate H 1 1 region physical properties and to disentangle the relative superposition of thermal and nonthermal sources (Kassim & Weiler 1990). We note the discrete region of extended nonthermal emission, coincident with the central molecular zone (CMZ), identified on Brogan’s 74 MHz image. Delineating its properties, including its spectrum, is challenging from observations at only one frequency. LOFAR’s frequency and resolution flexibility will allow its properties to be far better constrained. Its understanding could provide useful insights on the filling factor of hot gas and the magnetic energy density within the CMZ. 6.2 Cosmic Rays Kassim (1990) has described how absorption measurements of H I I regions at known distances probes the radial (3D) distribution and spectrum of Galactic cosmic ray clectrons. This constraint is akin to H I velocity measurements in determining the distribution of H 1 in the Galaxy. Combined with measurements of soft gamma-rays generated by collisions of thc synchrotron emitting electrons with matter, the strength and distribution of the magnetic field can also be extracted (Longair 1990; Webber 1990; Duric 2000). The aim is to link the high energy particles to the presumed acceleration sites in SNRs, addressing a key problem of Galactic cosmic ray origin. More than a dozen absorption regions have been detected at 74 MHz with the VLA, but LOFAR will increase the number of such measurements to hundreds (though not all of these will be in the GC). Frequency flexibility and resolution are crucial. The VLA is limited to a single frequency where the background emission is weak, and its limited surface brightness sensitivity renders only the largest H i I regions visible in absorption. For example, for LOFAR an H I I region of size 10’ at 30 MHz will produce an absorption depth of approximately a few Janskeys. However, for a more distant or smaller H 11 region only 1’ in size, the absorption depth is roughly 20 mJy. The 74 MHz VLA can detect only the largest such absorption regions while LOFAR should have the sensitivity to probe a range of sizes, even smaller than l’,thereby encompassing a much greater number of H 1 1 regions. 6.3
Large Scale Structure
Existing low frequency images (LaRosa et al. 2001) suggest that extensions of the Radio Arc and Sgr C appear to connect to larger scale structure seen on single dish maps, but the image fidelity falls short of quantifying the connection. Sofue (2000) has linked thesc extensions to large scale X-ray structure interpreted as energetic outflows driven by episodic bursts of star formation. LOFAR will provide surface brightness sensitivity capable of connecting structure seen on interferometer maps to those on single dish
N. E. Kassim et al.: LOFAR and the Galactic Center
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radio and large scale X-ray maps, addressing the possible link. This could be important in tracing manifestations of star formation possibly linked to previous episodes of black hole activity. The latter would be impossible to study directly because of Sgr A*’s current, presumably quiescent state.
7 Recombination lines
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Meter wavelength (- 120 MHz), stimulated recombination lines probe envelopes of extended H 11regions, delineating the distribution of hot (T lo4 K), low density (n 1cm-3) ionized gas. In the 20-100 MHz range, both hydrogen and carbon recombination lines occur from high quantum number transitions, in 100 cmp3) gas. At such high quantum numbers, atoms have radii cooler (T N 20-100 K), denser (7z of order 0.1 mm and are extremely sensitive to their surroundings, making them excellent probes of the ambient physical conditions of a cool, poorly delineated ISM component. Erickson et al. ( I 995) have mapped GC carbon line emission, which unexpectedly appears to trace 5 11 keV positron annihilation-line emission (Purcell et al. 1997; Erickson 2002). However, it is difficult to make a definitive association based on the current, single dish radio measurements. LOFAR could map this cool gas far more extensively, and investigate any correspondencc with the gamma ray emission. Physically, an association between positron annihilation emission and carbon line absorption is not predicted but is not unreasonable. The carbon lines trace extensive regions of fairly high density, cool gas which could be positron source regions via radioactive decay or regions where fast positrons thermalize and annihilate. If an association between the radio carbon lines and the gamma-ray emission could be established, it could provide kinematic distances and detailed physical information lacking from the 5 1 1 keV observations alone (Erickson 2002). N
8 Summary We anticipate that LOFAR will make important contributions to many areas of GC research. It may provide the lowest frequency detection of Sgr A*, improve our understanding of Sgr A West and East, and likely detect many new SNRs, PSRs, and NTFs. It will be able lo map the relative radial distribution of discrete sources as well as the spectrum and distribution of cosmic ray electrons towards the GC. Combined with gamma-ray measurements the latter may offer unique understanding of the origin of cosmic rays and the strength and distribution of the magnetic field towards the GC. LOFAR will detect many new small diameter sources and utilize them to map the distribution of ionized gas through separation of frequency dependent scattering and absorption effects. These will be complimented by recombination lines from low density, hot gas, and from lower frequency lines from cooler regions of hydrogen and carbon. Lastly, one of LOFAR’s key applications is as a transient detector, and it will provide the first sensitive and unbiased survey for GC radio transients. Its all-sky field of view, broad-band sensitivity, and ability to probe a continuous range of time scales for at least several hours each day, make it ideal for uncovering a hidden transient source population which current radio observations would never have seen. If recent X-ray observations serve as a proxy, there are indications of a potentially massive transient population awaiting discovery by LOFAR. Acknowledgements Basic research in radio astronomy at the NRL is supported by the Office of Naval Research. This research is supported at Sweet Briar College by Research Corporation and the Jeffress Memorial Trust. The National Radio Astronomy Observatory is a fdcility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.
References Beckert, T., et al. 1996, A&A, 307, 450 Bhatnagar, S. 2002, MNRAS, 332. 1 Brogan, C., el al. 2003, these proceedings
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Cordes, J. M. & Lazio, T. J. W. 1997, ApJ, 47.5. 557 Duric, N. 2000, in: Radio Astronomy at Long Wavelengths, R.G. Stone, K.W. Weiler, M.L. Goldstein, and J.-L. Bougeret (eds.), Gcophysical Monograph 1 19 (American Gcophysical Union, Washington, DC), p. 277 Erickson, W.C., et al. 1995, ApJ, 4.54, 125 Erickson, W.C. 2002, private communication Fender, R. P. & Kuulkers, E. 2001, MNRAS, ApJ, 324, 923 Gray. A. D. 1994a, MNRAS, 270, 861 Gray, A. D. 1994h, MNRAS, 270, 847 Green. D. A. 1991, PASP, 103,209 Haynes, R. F., et al. 1978, Aus. J. Phys. Astr., 45, 1 Hyrnan, S. D., ct al. 2002, AJ, 123, 1497 Hyman, S. D., et al. 2003, these proceedings Muno, M. P. 2003, these roceedin s Kassim, N. E., et al. l 9 8 g Nature,!322, 522 Kassim, N. E., et al. 1989, ApJ, 338, 152 Kassim, N. E. & Weiler, K. W. 1990, ApJ, 360, 184 KassIm, N. E., et a]. 1993, AJ, 106, 2218 Kassim, N. E. & Frail, D.A. 1996, MNRAS, 283, L51 Kassim, N. E. et al. 1999, in "The Central Parsecs of the Gaaxy", ASP Conference Series, Vol. 186, edited by H. Falcke, A. Cotera, W. J. Duschl, F. Melia, and M. J. Rieke, p. 403 Lang et al. 1999, AJ, 521, L41 LaRosa, T. N. & Kassim, N. E. 1985, ApJ, 299, L13 LaRosa, T. N., et al. 2000, AJ, 119, 207 LaRosa, T. N., Lazio, T. J. W., & Kassim, N. E. 2001, ApJ, 563, L163 Nord, M., et al. 2003, these proceedings Predehl, P., et al. 2003, these proceedings Purcell, W. R., el al. 1997, ApJ. 491, 725 Roy. S. & Pramesh, A. R. 2002, MNRAS, 329, 775 Shore, S. & LaRosa, T. N. 1999, ApJ, 521, 587 Sidoli, L., et al. 2001, A&A, 368, 835 Skinner, G. K., 1993, A&ASuppl., 97. 149 Sofue, Y. 2000, ApJ, 540, 224 Subrahmanyan, R. & Goss, W. M. 1996, MNRAS, 281,239 Yuan et al. 2003, these proceedings
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Astron. Nachr./AN 324, No. S1, 73 -77 (2003) / DO1 10.1002/asna.200385027
43 GHz SiO masers in late-type stars with 86 GHz SiO masers and astrometry with VERA in the Galactic center
'
Lorant 0. Sjouwerman* ,Maria Messineo', and Harm J. Habing2
' National Radio Astronomy Observatory, P. 0. Box 0, Socorro, NM 87801;
* Leiden Observatory P. 0. Box 9513. 2300 RA Leiden, The Netherlands;
Key words Masers, Galactic center, AGB stars, Circumstellar matter, Astrometry
We present Very Large Array (VLA) observations of 43 GHz SiO ( J = 1 --t 0 , v = 0 and v = 1) maser emission in a sub-sample of late-type stars in the Inner Galaxy. This sample of stars has relatively strong (> 0.5 Jy) 86 GHz SiO masers, and is located well within 2 degrees of the Galactic center (see contribution of Messineo et al. in these proceedings). It is our main intent to study the different maser properties of the sample, but as almost all 86 GHz SiO inasing stars are detected in at least one of the 43 GHz SiO maser lines, we here suggest that our selection criteria are a very efficient way to discover circumstellar 43 GHz SiO masers that can be used as gravitational mass probes. The 43 GHz SiO masers are extremely useful for VLBI astrometric and proper motion measurement\, for example with the Japanese "VLBI Exploration of Radio Astrometry" (VERA) VLBT network, allowing determination of Galactic structure, geometric distances (to individual objects, but also e.g. to the Galactic center), and types o f orbits of individual stars, in particular in the central part of the Galaxy.
1 Introduction As shown by Lindqvist et al. (IY92a), OH/IR stars are very useful objects to probe the gravitational potential in the central part of the Galaxy. OH/IR stars are the more massive Asymptotic Giant Branch (AGB) stars, that at the tip of the Giant branch are variable and loose mass at a very high rate due to Hydrogen shell burning and Helium flashes (thermal pulses). Their circumstellar envelope totally obscures the central star in the visible, re-radiates the stellar radiation in the mid-infrared. and - if proper conditions are met frequently sustains maser emission at 1612 MHz from the OH-molecule (e.g. Habing 1996). Circumventing the high visible extinction i n the Galactic center region, the OHllR stars are very prominent in the mid-infrared and radio wavelengths. With some effort the bolometric magnitudes can be measured (Blommaert et al. 1998; Wood et al. 1998; Ortiz et al. 2002) and their masers reveal their line-of-sight velocities accurate to a fraction of one km s-' instantaneously (e.g. Baud et al. 1981). About 200 OH/IR stars are known in the central degree of the Galactic center (Lindqvist et al. 19Y2b; Sjouwerman et al. IYYXb and references therein), and a few hundred more in the Inner Galaxy. In this paper we refer to the Inner Galaxy as Galactic longitudes 1, with 30" < I < -30" (e.g. Sevenster et al. 1997). Because of severe interstellar scattering, which scales as A', in the Galactic center the 1612 MHz OH masers arc not useful for astrometry (e.g. van Langevelde et al. 1992). OWIR stars frequently also harbor 22 GHz H a 0 andlor 4 3 GHz SiO masers in their circumstellar shell (e.g. Habing 1996 and refercnccs therein). While the HzO masers are highly variable, the 4 3 GHz ( J = 1 + 0) SiO maser is more stable. Thus in particular the 43 GHz S i O maser can be used for astrometry, most notably in the Galactic center region (Menten et al. 1997; Sjouwerman et al. 1998a, 2002; Reid et al. 2003). However, the number of43 * Corresponding author: e-mail:
lsjouwermanQnrao.edu @ 2003 WILEY-VCH Vcrlng GmhH B Co KGaA. Wr~ahem
14
L. 0. Siouwerman et al.: 43 GHz SiO masers and astrometrv with VERA in the GC
GHz masers in OH/IR stars in the Galactic center is only a very small fraction of the known OH masers (Sjouwerman 1997; Sjouwerman et al. 2002), which has triggered searches for 43 GHz masers in other types of mid-infrared sources and in blind surveys in the Galactic center (Menten et al. 1997; Sjouwerman 1997, and the Japanese groups as published in Shiki et al. 1998; Izumiura et al. 1998; Miyazaki et al. 2001; Deguchi et al. 2002; Imai et al. 2002). In particular, it is believed that the OWIR stars arc just the tip of the iceberg, and that the 43 GHz maser, or even the 86 GHz ( J = 2 + I , 'u = 1) SiO maser, is readily observable in the much more numerous Mira stars, and in some Semi-regular AGB variables and red supergiant stars (Habing 1996). The number of 43 GHz SiO masers available for astrometric studies, kinematics, dynamics and proper motions in the Galactic center has been increasing - slowly, but steadily, and recently very rapidly (Imai et al. 2002; Sjouwerman et al. 2002 and references therein; this work).
2 86 GHz SiO masers Unlike the 43 GHz SiO maser surveys of the Japanese groups mentioned above, our group has chosen to focus on the 86 GHz SiO maser to obtain more stellar mass probes to investigate the structure of the Galaxy (Messineo et al. 2002, and Messineo et al. i n these proceedings). The 86 GHz masing stars in our sample have a lower mass-loss rate than the OH/IR stars. Our survey is therefore complementary to previous OH maser surveys, although we trace-a slightly different population (e.g. Habing 1996). The combination of the OWIR stars with the population of stars defined through their 86 GHz SiO maser, should allow a proper study of the kinematics and structure of the Inner Galaxy. Because a targeted survey is more efficient than a blind survey (Messineo et al. 2002), we selected mid-infrared sources with colors and magnitudes characteristic for Mira-like stars from the ISOGAL (Omont ct al. 1999) and MSX (Egan et al. 1999) catalogs. That is, characteristic for dense circumstellar envelopes that could harbor masers, but avoiding the highest mass-loss OH/IR stars. The results of the 86 GHz SiO maser survey are given in this volume by Messineo et a]. (see also Messineo et al. 2002).
3 43 GHz observations of 86 GHz SiO masers Apart from studying Galactic structure, we are also interested in what determines whether an AGB star has a particular maser. It is possible that, among other factors, the mass-loss rate has a major influence on the type of maser apparent in the circumstellar environment. The high mass-loss rate (- lop4 Mayr-l) OWTR stars may be too turbulent close to the star to sustain 86 GHz masers, whereas semi-regular (SR) low mass-loss rate (Moyrpl) stars may never build up a dense enough shell at large distance to form an OH maser. For this reason, we arc conducting a more general search for the different masers in the different types of late-type AGB stars. Here we report on using the Very Large Array (VLA) to target the brighter 86 GHz masers found by Messineo et al. (2002) for 43 GHz (.I = 1 + 0, u = 0 and v = 1) SiO masers in January 2002; we are currently continuing our search in the less bright 86 GHz masers. Because of a limited relative bandwidth at 43 GHz, the VLA is not very effective in searching for 43 GHz maser emission ranging from -400 to +400 km spl in the ISOGAL and MSX sources. However, because we have a-priori information on the stellar velocity from the 86 GHz maser detections, the VLA, with its large collecting area, is a very powerful instrument to detect the corresponding 43 GHz masers within a few minutes of observation. The only limitation is imposed by the rapid phase changes over the array, requiring a phase calibrator within about two degrees of the targeted source. Fortunately, Sgr A* itself, the black hole in the dynamical center of the Galaxy located at the highest apparent stellar density, happens to be a very good calibrator. We searched 39 of our 86 GHz SiO masers and found 43 GHz SiO masers in 38 of them. Typical flux densities range from 0.1 to 2 Jy, which makes them observable with the VLBI Exploration of Radio Astrometry (VERA) network. Although we targeted the stronger 86 GHz masers in our first 43 GHz maser search, because of the high detection rate (38/39) it is likely that a large fraction of the less strong 86 GHz
75
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Fig. 1 The 86 GHz SiO masers from Messineo ct al. (2002; big dots) compared to the location of the OH sources of both Lindqvist et al. (1992b; crosses), Sjouwerman et al. (199%; also crosses) and Sevenster et al. (1997; triangles). This figure shows that the 86 GHz SiO masers fill the gap in kinematic data between the deep, but small sky-coverage OH surveys of Lindqvist et al. (199%) and Sjouwernian e~ al. (1998b), and the large but less sensitive OH sinvey of Sevenster et al. (1997) covering the full Inner Galaxy. Most 43 GHz SiO masers within about one degree of the Galactic center found to date form a subset of the 86 GHz SiO masers (this work) and known Long-Period variables -only a few have been found in blind surveys (references in the text).
0
-1
Galactic Longitude (degrees)
Galactic Longitude (degrees)
Fig. 2 Current status of the 43 GHz SiO masers in the Galactic center. The top (square) panel shows a blow-up of the central region of the lower panel, and symbols used are indicated by the name of the PI reporting the 43 GHz SiO maser detection (although Messineo et al. 2002 refers to this 43 GHz work). Note that no ISOGAL sources are located close to the very center, because the observations avoided the most luminous infrared sources to protect the detector and save on cooling fluid. However, this niche has nicely been filled by Imai et al. (2002) who have targeted the Long-Period variables identified after a few year's monitoring program by Glass et al. (2001). A few detections are not shown because only a crude position is known (e.g. Izumiura et al. 1998), or because they blend in in the central region (e.g. Menten et al. 1997). The 43 GHz masers found here extend thc possible use of 43 GHz masers out from the inner 30 pc to more than 60 pc from the Galactic center without the need of long near-infrared monitoring programs (Glass et al. 2001).
masers in o u r follow-up V L A observations will also show 43 GHz S i O maser emission . This makes a targeted survey for 43 GHz masers in stars with 86 G H z masers, selected using their mid-infrared colors, a very efficient way to find 43 G H z masers for follow-up VLBI observations. Here w e suggest that an efficient way to discover 43 G H z S i O masers is t o target the 86 G H z SiO masers found in the sources selected from the ISOGAL/MSX-DENIS color-magnitude diagram (Messineo et al. 2002).
4
Astrometry with 43 GHz SiO masers
Once the position of an AGB star is known accurately, say to an arcsecond or better, the star becomes very attractive for astrometry purposes. Menten et al. (1997) have demonstrated nicely how the infrared and radio reference frames in the Galactic center can b e aligned by using AGB and supergiant stars; the only objects that radiate both in the infrared and in the radio through their obscuring circumstellar shell and
L. 0. Sjouwennan et al.: 43 GHz SiO masers and astrometry with VERA in the GC
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masers therein. Because the position and the line-of-sight velocity of the AGB stars can be measured very accurately using the 43 GHz SiO maser (e.g. Sjouwerman et al. 1998a), it should be possible to obtain information on the individual orbits of the stars (Reid et al. 2003). This orbital information is important in deriving the three dimensional characteristics of the gravitational potential in the Galactic center. As mentioned above, in the Galactic center the 1612 MHz OH maser is too scattered to be used for accurate astrometry (e.g. van Langevelde et al. 1992). Not only is the scattered OH maser too extended (500 mas) to be detected with VLBI, but also the intrinsic position of the OH maser cannot be determined better than a few 100 mas. For the Galactic center, a typical line-of-sight velocity of 100 km s-l translates to 2.6 mas yr-', i.e. too small relative to the intrinsic position uncertainty to be measured accurately in a few year's time. At higher frequencies, the interstellar scattering becomes less dominant and much more accurate positions can be determined. For example, the intrinsic and scattered size of the circumstellar 22 GHz HzO maser is about 2 mas. But the H2O maser is probably too variable to be a good candidate for position monitoring. With an one mas intrinsic (and scatter) size, the 43 GHz SiO maser is the most promising candidate for determining accurate stellar positions in the Galactic center. Although a position of an 86 GHz SiO maser would be completely intrinsic, with a negligible scatter size (to a fraction of a mas), the current instruments and weather in the upper atmosphere seem to hinder mas-accuratc position measurements in the Galactic center.
5
VLBA observations of 43 GHz SiO masers
Driven by the prospect of measuring stellar proper motions at 43 GHz, Sjouwerman et al. (1998a) have used the Very Large Baseline Array (VLBA) and phased VLA to obtain first epoch mas-accurate stellar positions using the circumstellar 43 GHz SiO maser. See also Reid et al. (2003) for a more recent result. To accommodate for the rapid tropospheric phase fluctuations caused by the troposphere before fast-switching at the VLA was implemented, Sjouwerman et al. (1998a) divided the VLA up into 2 subarrays - one observing the targeted 43 GHz maser, the other observing the phase calibrator - while the VLBA was phase-referencing with a 20/20 second cycle. Although only 2 mas-accurate positions out of 10 sources were measured, the feasibility was demonstrated (Fig. 3). Since then the available instruments
!.I
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$; j
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Fig. 3 Two detections in the 43 GHz SiO maser VLBA astrometry campaign (taken from Sjouwerman et al. 1998a).
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have improved; in particular the dual beam VERA will eliminate the need for fragile phase-referencing techniques.
6 Prospects with the VERA network With projected baseline lengths of up to approximately 1700 km, the VERA network will have a resolution ofjust slightly less than one mas at a frequency of43 GHz. This is exactly the proper resolution to observe circumstellar 43 GHz SiO masers in the Galactic center! Most of the 43 GHz SiO masers in the Galactic center found to date are located within 15 arcminutes, or 30 parsecs in projection, of the dynamical center, the black hole Sgr A*, and are concentrated toward the stellar density maximum. However, this (still continuing) work has also located many additional circunistellar 43 GHz SiO masers out to about 2 degrees (300 pc) or the Galactic center. Typical Hux densities of the circumstellar 43 GHz SiO masers are several 100 mly, and about 4 krn s p l wide. It is fortunate that Sgr A*, point-like and with a continuum Hux density of about 1-2 Jy at 43 GHz, is a good calibrator and located in the middle of all circumstellar 43 GHz SiO masers in the Galactic center. As Sgr A* is the dynamical center, relative astrometry with respect to Sgr A* is absolute astrometry with respect to the dynamical center of the Galaxy! The sensitivity of the VERA network, with four 25 meter antennas, will at first be a bit low to detect the majority of the circumstellar 43 GHz SiO masers i n the Galactic center. However, utilizing the dual beam receivers, spanning an angle of 2 degrees between the calibrator (Sgr A*) and the target source, all the currently known circumstellar 43 GHz SiO masers in the Galactic center can be observed coherently with long integration times. Milli-arcsecond accurate astrornetric positions can be measured with VERA and will provide proper motions of individual stars i n the Galactic center, and possibly even individual stellar orbits in the inner few tens of parsecs of the Galactic center in the next decade.
References Baud, B., Habing, H. J., Matthews, H. E., Winnherg, A . 1981, A&A 95, 156 Blommaert, J. A. D. L., van der k e n , W. E. C. J., van Langevelde, H. J . , Habing, H. J.. Sjouwerman. L. 0.. 1998, A&A 339,991 Deguchi, S., Fujii, T., Miyoshi, M., Nakashiina. J. 2002, PASJ 54, 61 Egdn 1999, A S P Conf. Ser. 177,404 Glass, I S . , Matsumoto, S., Carter, B. S., Sekiguchi, K . 2001, MNRAS 321, 77 Habing, H. J. 1996, A&A Rev 7,97 lmai, H., Deguchi, S., Fujii, T., Glass, 1. S., Ita, Y., Izumiura, H., Kameya, O., Miyazaki, A,, Nakada, Y., Nakashima, J. 2002, PASJ Let 54, L19 Izumiura, H., Deguchi, S., Fujii, T. 1998, ApJ Let 494, L89 Lindqvist. M., Habing, H. J., Winnberg. A. 19928, AXrA 259, 118 Lindqvist. M., Winnberg. A., Habing, H. J., Matthews, H. E. l992b, A & A Sup 92,43 Menten, K. M., Reid, M. J., Eckart, A,, Genzel, R. 1997, ApJ Let 475, LI 1 1 Messineo, M., Habing, H. J., Sjouwerman, L. O., Omont, A,, Menten, K. M. 2002, A&A 393, 115 Miyazaki, A,, Deguchi, S., Tsuboi, M., Kasuga, T., Takano, S. 2001, PASJ 53, 501 Omonl, A., GdneSb, S., Alard, C., Blommaert, J. A. D. L., Caillaud, B., Copet, E., Fouque, P., Gilmore, G., Ojha, D., Schultheis, M., et al. 1999, A&A 348, 755 Ortiz, R., Blommaert, J. A. D. L., Copet, E.. Ganesh, S., Habing, H. J., Messineo, M., Omont, A., Schultheis, M., Schuller, F. 2002, A&A 388, 279 Reid, M. J., Menten, K. M., Genzel, R., Ott, T., Schiidel, R., & Eckart. A. 2003, ApJ in press, (aslro-ph 0212273) Sevcnster, M. N., Chapman, J. M., Habing, H. J., Killeen, N. E. B.. Lindqvist, M. 1997, A&AS 122, 79 Shiki, S., Ohishi, M., Deguchi, S. 1998, ApJ 487, 206 Sjouwerman, L. O., 1997, PhD thesis Onsala, Sweden Sjouwerman, L. 0.. Lindqvist, M., van Langevelde, H. J., Diamond, P. J. 2002, A&A 391, 967 S.jouwerman, L. O., van Langevelde, H. J., Diamond. P. J. 1998a. A&A 339, 897 Sjouwerman, L. O., van Langevelde, H. J., Winnberg, A,, Habing, H. J. 1998b, A&A Sup 128, 35 van Langevelde, H. J., Frail, D. A,, Cordes, J. M., Diamond, P. 5. 1992, ApJ 396, 686 Wood, P. R., Hahing, H. J., McGreggor, P. J. 1998, A&A 336, 925
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A\tron. Nachr./AN 324, No. S I , 79 - 83 (2003) / DO1 10.1002/asna.200385028
A Search for Radio Transients at 0.33 GHz in the GC Scott D. Hyman*',T. Joseph W. Lazio', Namir E. Kassim2, Michael E. Nard'.', and Jennifer L. Neureuther'
' Department of Physics, Sweet Briar College. Sweet Briar, VA 24595 USA
' Naval Research Laboratory. Code 72 13, Washington.DC 20375-5351 USA University of New Mexico. Department of Physics and Astronomy, 800 Yale Blvd. NE, Alhuquerque, NM 87131 USA
Key words radio transients. low radio frequencies, Galactic center PACS 04A25
We report on a search for transient and variable radio sources in the Galactic center using a number of 327 MHz VLA observations made during the 1990's, and a series of monthly VLA observations made during Spring and Summer 2002. A typical yield of compact sources in a given epoch is roughly 200. We have detected one new bright radio transient, GCRT 51746-2757, located only 1.1 degrees north of the Galactic center. We discuss our on-going transient monitoring program and the implications of'this work for constraining the Galactic center population of transients.
1 Introduction Known classes of highly variable and transient radio sources include radio counterparts of X-ray sources and microquasars. Although there are many examples of variable radio sources discovered as a result of high-cncrgy observations, there are surprisingly few radio surveys for highly variable or transient sources. A radio survey of the Galactic plane (Gregory & Taylor 1981, 1986) discovered 4 variable sources including GT 0236+610, a Galactic X-ray binary, and 1 candidate transient. The MIT-Green Bank surveys (Langston et al. 1990; Griffith et al. 1990; Griffith et al. 1991) discovered a number oT variable sources (< 40% variable). An on-going program at NRAO Green Bank monitors the Galactic planc at 8.4 and 14.4 GHz (Langston et al. 2000). The Galactic center (GC) is a promising region in which to search for highly variable and transient sources. The stellar densities are high, and neutron star and black hole binaries appear as (transient or variable) X-ray sources concentrated toward the GC (Skinner 1993). Previous surveys have been ill-suited for detecting radio transients toward the GC, however. Typically they have utilized either single dish instruments. which suffer from confusion in the inner Galaxy, or they have utilized the VLA for only a single epoch (e.g., Zoonematkermani et al. 1990; Becker et al. 1994). The first two radio transients detected toward the GC were A1742-28 (Davies et al. 1976) and the Galactic Center Transient (GCT, Zhao et al. 1992). These two transients had similar radio properties, but only the former was associated with an X-ray source. More recently, radio counterparts to the Xray transients, XTE 51748288 (Hellming et al. 1998a, 199%; Rupen et al. 1998) and GRS 1739278 (Hellming et al. 1996) have been detected in the GC. The G C also contains many exotic phenomcna not seen elsewhere in the Galaxy and may contain additional previously undetected novel classes of radio sources. This paper reports on the current status of our transient monitoring program utilizing low-frequency VLA radio observations. We summarize the detection and properties of a new, bright f-200 mJy) radio * Corresponding author: c-mail:
shyrnanQsbc.edu. Phone: + I 434381 6158, Fax: + I 434381 6488 @ 2003 W I L R Y ~ V C HVerlag GmhH B Cn KCiaA. Wcinhcim
S. D. Hvman et al.: Radio Transients
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transient, GCRT 51746-2757 (Hyman et al. 2002), and the low-frequency detection of the radio counterpart to the X-ray transient, XTE 51748-288. We present one notable variable radio source candidate, and describe efficient, automated techniques we are developing to search for additional transient and variable sources. We also discuss the impact of our results on our long-term goal to constrain the nature of the transient and variable source population(s) based on their individual and group properties, and outline our future monitoring plans.
2 Radio Observations In Spring and Summer 2002, we initiated a high resolution, high sensitivity monitoring program of the GC with the VLA at 0.33 GHz in order to detect highly variable and transient radio sources. Recent developments in low frequency 3-dimensional imaging techniques (LaRosa et al. 2000; Nord et al. 2003a) allow us to produce wide-field images (= 2") with uniform and high resolution across the field. The large field-of-view covers the entire G C region with a single observation. Low frequency observations also increase the likelihood of detecting transient sources since they typically have steep spectra. We have used the 2002 observations, along with archival VLA observations made during the 1990's, to search for variable and transient candidates. We have established temporal baseline measurements for -200 sources and detected two bright (-200 mJy) radio transients (Hyman et al. 2002) and a few variable candidates. Table 1 describes our database of A- and B-configuration VLA observations consisting of various integration times, resolutions, and sensitivities. Table 1 VLA 0.33 GHz Observations Epoch 1989 March 1995 August 1996 October 1997 February 1998 March 1998 September 2002 March 2002 April 2002 May 2002 June 2002 July 2002 combined
VLA Resolution
20" x 40" 5" x 10" 51' x 10" 10" x 20" 51' x 10" 20/' x 40" 5" x 10" 5" x 10" 10" x 20" 15" X 60" 15" x 40" 8" x 14"
Duration (hr) 5.5 1.0 5.x 1.3 5.5 6.7 1.1 I.4 I .4
0.6 I .0 5.5
rmy
(mJy bm- ') 5 II 3 14 3 3
7 7 5 I4 9 3
While the inhomogeneity of the observing parameters listed in Table 1 has hindered reliable detection of variability for fainter sources, the observations have allowed us to provide constraints on the timescales of brighter transient and variable sources. For example, Figure 1 depicts the variable flux density detection threshold (70) for sources varying during our 1998 March observation. Simulated transient sources with timescales, T, shorter than the total integration time, T = 5.5 hr, were added to the visibility data. After reimaging, the detection threshold was verified to increase by a factor of T/r, as indicated by the negatively sloped line in the figure. Note that this short timescale threshold would be lower still by (T/T)'/' for images synthesized from a subset of the observation time equal to T. We are presently developing an iterative, automated routine to accomplish this, in order to improve our sensitivity to extremely short timescale transients. Figure 1 also shows the constant 7 u threshold curve for sources varying with timescales, T > T, although presently our database includes only a few observations with sensitivities similar to that of 1998 March with which to sample this temporal regime. Constant brightness temperature curves are also indicated based on source size upper limits (= CT). The temperature upper limit for an incoherent synchrotron emitter is 10l2K.
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Fig. 1 Threshold for transient detection as a function of timescale, T .We are sensitive to detecting sources with flux
densities in the region above the solid lines. The four parallel dashed lines are constant brightness temperature curves (left to right: 10"K, 10l°K, 10'K. and 106K) based on source size upper limits (= CT).
3 Search Methodologies Various methods were employed to search for transient and variable sources on each or the epoch images. First, the CLEAN components of the 5.5 hr 1998 A- and B-configuration images were subtracted from the visibility data obtained in similar configurations for each of the other epochs using the task UVSUB in AIPS. Wide-field imaging requires that many small, overlapping facets be imaged across the field-ofview in order to prevent significant distortion of sources (LaRosa et al. 2000). Accordingly, the residual data were imaged using 512 ( 5 5 ) facets for the higher (lower) resolution epochs. Each facet was then searched for transients identifiable either as a bright source, indicative of newly detected source, or as a "hole" in the image, indicative of a source no longer detectable. This "model subtraction" technique mitigates confusion, removes the time-consuming and computation-intensive CLEANing step, and results in superior sensitivity than would be obtained otherwise for epochs with shorter integration times. For epochs of similar integration time (e.g., 1996 October), the removal of the deconvolution (CLEAN) step is still a clear advantage. Figure 2 shows the bright transient, GCRT 51746.2757, visible only in the 1998 September image (see Sec. 4), appearing as a "hole" source in the 1989 March residual image. The other sources in the field are largely removed by the model subtraction technique, although not completely, due to incomplete CLEANing of the 1998 model image. Because of this limitation and other systematic uncertainties leading to imaging artifacts, the model subtraction technique is adequate only for detecting transients or highly variable sources. Therefore, only sources exceeding a threshold peak flux density of 7 ~ 7were designated as transient or variable candidates by this mcthod. In the second detection method, used primarily to search for variable sources, model subtraction was not employed. Instead, an image was synthesized for each epoch and the automated source detection program SAD (AIPS task Search and Destroy) was run. We compared source measurements with those listed in our reference database. The latter consists of -250 sources detected on an image synthesized from the combination of our 1996 and 1998 observations (Nord et al. 2003b). This comparison enabled us to quickly remove false detections resulting from the SAD procedure. Conversely, sources reliably detected on the superior sensitivity combined image, lend credence to marginal detections made on individual epoch images. All sources varying by more than 50 between any two epochs were remeasured, and if confirmcd, these sources were designated as variable candidates. Sources that differed due to unmatched resolutions and/or confusion with extended emission were removed from the list. Only a few variable candidates were
82
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Fig. 2 Images of GCRT 51746-2757 at 0.33 GHz. (Lefr) 1989 March residual image after subtraction of 1998 September visibility data. The radio transient appears as a “hole” in the image. (Right) 1998 September detection.
Fig. 3 Images of the variable candidate G0.490-1.043 at 0.33 GHL. (Lefr) No detection on 1989 March 18 image with rms senstivity of 7 mJy beam-’. (Righf)80 mJy detection on 1998 September 25 image.
detected. One is G0.490- 1.043, which has consistent -80 mJy detections except for a non-detection in 1989 March, as shown in Figure 3.
4 Confirmed Transients Our monitoring program has detected two transients, GCRT 51746-2757 and XTE 51748-288, both in the 1998 September image, and the latter also in a lower resolution 0.33 GHz image made from 1998 Novembcr observations. Details of these detections are presented in Hyman ct al. (2002). In summary, GCRT 51746-2757 is located only 1.1” (150 pc in projection) north of Sgr A* with a flux density at detection of 21 6 f 20 mJy. Non-detections in follow-up VLA observations at higher frequencies in 2000 July and December indicate that the source faded significantly in the interim. While GCRT 51746-2757 was detected at only one epoch and at only one frequency, a number of diagnostic imaging tests were done which provide evidence that the detection is robust. Contemporaneous searches revealed no X-ray counterpart. We conclude that GCRT 51746-2757 was either highly Doppler-boosted in the radio or was a “fast” X-ray transient, or that it is a member of a class of radio transients with no associated X-ray emission. The radio counterpart to the X-ray transient, XTE J1748-288, has been monitored extensively at 1.5 GHz and higher frequencies (Hjellming et al. 1998~1,1998b; Rupen et al. 1998). Our 0.33 GHz detections
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(-250 mJy) in 1998 September and November together with those at higher frequencies (kindly provided by M. Rupen) indicate that the spectrum of X T E 51748-288 steepened significantly over a period of months after its radio peak in 1998 June. Based o n the ratio of its radio and X-ray fluxes, Fender &L Kuulkers (2001) have classified X T E 51748-288 as a black-hole binary candidate. Of the -30 Low-Mass X-ray Binary sources reported within our 0.33 GHz field-of-view, w e have only detected a radio transient counterpart t o X T E 51748-288. We have detected the radio counterpart to the X-ray transient G R S 1734-292, however, which underwent a burst in 1992 September. Based on its redshift, Marti et al. ( 1 998) classified this source as a Seyfert 1 galaxy. They found n o radio variability at 1 .5 GHz and higher frequencies, consistent with our nearly constant measured flux density of -150 mJy at 0.33 GHz. Evidently G R S 1734-292 is presently in a quiescent state.
5
Conclusions
V L A observations of the Galactic center at 0.33 GHz provide high-resolution, high dynamic range, large lield-of-view radio images which are well-suitcd for transient monitoring. From recent and archival 0.33 G H z observations, w e conclude that radio transients above -50 mJy are either very infrequent (approximately one every few years) o r have timescales much shorter than a month. T h u s far, w e have found only a few radio variable candidates out of -250 sources detected, but continued monitoring may yet reveal additional variables. The task is more difficult than detecting transients; e.g., a 50 mJy source visible in only one epoch is a more reliable detection than a 50 mJy variation in a 300 mJy source. T h e latter could b e due to a number of systematic uncertainties. The scarcity of radio transients/variables detected in this survey so far, underscores the need to initiate a far more extensive monitoring program. Increased sensitivity, and inore frequent observations, are required in order to detect, monitor, and identify a large number of radio sources transientlvariable over very short (< 1 hr) to very long (> 1 yr) time scales. Acknowledgements J.L.N. and S.D.H. thank Grant Denn and Manana Lazarova for further Large Array of the National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement with Associated Universities, Inc. This research was supported at NRL by the Office of Naval Research, and at SBC by Research Corporation and the JetTress Memorial Trust.
References Becker, R. H., White, R. L., Helfand, D. J., & Zoonematkermani, S. 1994, ApJ Suppl., 91, 347 Davies, R. D., Walsh, D.. Browne, I. W. A., Edwards, M. R., & Noble, R. G. 1976, Nature, 261,476 Fender, R. P., & Kuulkers, E. 2001, MNRAS, 324,923 Gregory, P. C . , &Taylor, A. R. 1986, Astron. .I.92, , 371 Gregory. P. C . , &Taylor, A. R. 1981, ApJ, 241. 596 Griflith, M., Hetlin, M., Conner, S., Burke, B., & Langston, G. 1991, ApJ Suppl., 75, 801 Griflith, M., Langston, G., Heflin, M., Conner, S., Lehhr, J., & Burke, B. 1990, ApJ Suppl., 74, 129 Hjellming, R. M., Rupen, M. P., & Mioduszewski, A. J. 1998a, IAUC 6934 Hjellming, R. M., Rupen. M. P., Ghigo, F., Waltnian, E. B., & Mioduszewski, A. J. 1998b, IAUC 6937 Hjellming, R. M., Rupen, M. P., Marti, J., Mirabel, F.. & Rodriguez, L. F. 1996, IAUC 6383 Hyman, S. D., Lazio, T. J. W., & Kassim, N. E. 2002, Astron. J., 123, 1497 Langston, G., Minter, A., D’Addario, L., Eberhardl, K., Koski, K., & Zuber, J. 2000, Astron. J., 1 19, 2801 Langston, G . I., Heflin. M. B., Conner, S. R.. LehBr. J., Carilli, C. L., & Burke, B. F. 1990, ApJ Suppl., 72, 621 LaRosa, T. N., Kassim, N. E., Lazio, T. J. W., & Hyman, S. D. 2000, Astron. J., 119, 207 Lazio, T. J . W., & Cordes, J. M. 1998, ApJ, 505, 715 Marti, J., Mirabel, 1. F., Chaty, S., & Rodriguez, L. F. 1998, A&A, 330, 72 Nord, M. E., et al. 2003a. these proceedings. Nord, M. E., et al. 2003b, in preparation. Rupen, M . P., Hjellming, R. M., & Mioduszewski, A. J . 1998, IAUC 6938. Skinner, G. K. 1993, A&A Suppl., 97, 149 Zhao, J-H., Roberts, D. A., Goss, W. M., Frail, D. A., Lo, K. Y., Subrahmanyan, R., Kesteven, M. J., Ekera, D. A,, Allen, D. A., Burton, M. G., & Spyromilio. J. 1992. Science, 255, 1538 Zoonernatkermani, S., Helfand, D. J., Becker, R. H., White, R. L., & Perley, R. A. 1990, ApJ Suppl., 74, 181
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Astron. NachrJAN 324. No. SI. 85-91 (2003)/ DO1 10.1002/aana.200385088
A Molecular Face-on View of the Galactic Center Region Tsuyoshi Sawada* I , Tetsuo Hasegawa2, Toshihiro Hands?, and R. J. Cohen'
' Nobeyama Radio Observatory
' National Astronomical Observatory of Japan ' Institute of Astronomy, University of Tokyo University of Manchestcr
Key words Galaxy: center, ISM: molecules
Abstract. We present a method to derive positions of molecular clouds along the lines of sight from a comparison between 2.6 mm CO emission lines and 18 crn OH absorption lines, and apply it to the central region of the Milky Way. With some simple but justifiable assumptions, we derive a face-on distribution of the CO brightness and corresponding radial velocity in the Galactic center without the help of kinematical models. The derived face-on distribution of the gas is elongated and inclined so that the Galactic-eastern (positive longitude) side is closer to us. The gas distribution is dominated by a barlike central condensation, whose size is about 500 x 200 pc. The major axis of the condensation is tilted with respect to the line of sight by an angle of ru 70" (tilled by = 10-50" from the large-scale stellar bar). This geometry resembles the central regions of barred galaxies. The velocity field shows highly noncircular motion in the central condensation. These characteristics agree with a picture in which the kinematics of the gas in the central hundreds of parsecs of the Galaxy is undcr the strong influence of a barred potential.
1 Introduction The behavior of molecular gas, in particular its physical conditions and kinematics, in central regions of galaxies is key information to understand the star forming activity which occurs there. The Galactic center can be observed in much greater detail compared with central regions of other galaxies. It has long heen argued that the Galaxy has a bar, and recent studies strongly suggest the existence of a bar, whose Galactic-eastern (positive longitude) end is closer to us (see Gerhard 1999, for reviews). Therefore the Galactic center is also important for thc study of phenomena in central regions of barred galaxies. However, its inevitable edge-on perspective sometimes complicates the interpretation of the data. In particular, a face-on image o f the Galactic center is very hard to construct, though such an image would be very helpful to understand its kinematics and to make a comparison with central regions o f other galaxies. Attempts have been made to construct models of gas kinematics (see, e.g., Liszt & Burton 1980; Binney et al. 1991). Kinematical models can be used to project position-velocity diagrams of molecular line data into a face-on view (see, e.g., Cohen & Dent 1983; Sofue 1995). This is an indirect method to investigate the spalial distribution of the gas: it would be invaluable if we could derive positions and motions of molecular clouds independent of kinematical assumptions. Wc present a method to derive a molccular face-on view o f the Galactic center without any help of kinematical models. In $ 2 we describe the hasic methodology. Using that, we draw a face-on distribution of the molecular gas from existing data and discuss the resultant face-on view in 3 3. * e-mail: sawadaQnro.nao.ac.jp
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2 TheMethod 2.1 Qualitative Inspection of Molecular Gas Distribution
The essence of our method lies in comparing molecular emission and absorption spectra. We compare the CO 2.6 iiiiii emission with the OH 18 crn absorption. Because the Galactic center region itself i s an intense diffuse 18 crii continuum source, strong OH absorption arises preferentially from the gas that lies in front of the continuum regions, rather than gas that lies behind them. On the other hand, the CO emission samples the gas both in front and back of the continuum sources equally. Thus the OH/CO ratio carries information on the position of the gas along the line of sight rclative to the continuum sources. We used the "CO J = 1 - 0 data by Bitran et al. (1997) and an absorption survey of the OH main lines (1665 MHz and 1667 MHz) made by Boyce & Cohen ( I 994). The CO data were smoothed and resampled with a 0?2 grid to match the OH data. Figure 1 shows the spectra of the OH absorption and the CO emission at (a. b) = (-002.000). For example, we may draw attention to the velocity components near U L S R E -130 kin s-' and 160 knis-' (shadowed in Fig. l), both of which belong to the so-called expanding molecular ring (EMR; Kaifu, Kato, & Iguchi 1972, Scoville 1972). The CO intensities of these components are almost the same: i.e., the amounts of molecular gas are similar. On the other hand, the OH absorption depths are strikingly different. We immediately deduce from this fact that the negative-velocity component, at which deep absorption i s seen, is located in front of strong continuum source surrounding the Galactic center, while the positive-velocity component is behind it. This logic led Kaifu, Kato, & Iguchi (1972) to conclude that this feature is expanding away from the center. Figure 2 shows the longitudc-velocity (l-w) diagram of the ratio between the OH apparent opacity ( T =~Tat,s/Tc,,,,t,; ~ ~ ~where Tabs and Tc,,,t are line absorption and continuum antenna temperatures, respectively) and the CO J = 1 - 0 line intensity. The ratio has a clear trend; high ratio tends to he seen at positive !, negative 21. The hatched velocity ranges in Fig. 2 are excluded from the following analysis because the ratio is affected by the foregroundgas in the Galactic disk. 300
I
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'
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'
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'
/ , I , I , I , l , I 1 1 , 1 ,I , 3 2 I 0 1 - 2 - 3 - 4 - 5 - 6 GaldCtlC
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Fig. 1 Sample spectra of CO emission (thin line) and OH absorption (thick line) at ( B , b) = (-0?2.0?0). Shadowed velocity components demonstrate the difference of cloud position along the line of sight (see text).
0 05
Longitude [degrees]
0 10
n 15
(OH Apparent Opacity)/(CO J = 1 4 Tim) [K '1
Fig. 2 The ratio between the OH apparent opacity and the CO J = 1 - 0 intensity at b = 0". Contours are TL,u(CO)= 1 , 2 , 4 , 7 , 1 0 ,and 14 [K]. Hatchedvelocity
ranges were excluded from the analysis because of contamination by clouds well outside the Galactic nucleus.
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Deriving Distances to Clouds
We extend the logic shown above to quantitative estimation of molecular gas distribution. In order to determine the distances to molecular clouds quantitatively, we adopt the following four assumptions. (1) At a given P,emission observed at each velocity bin comes from a single position along the line of sight. (2) The CO line intensity Tc.0 at a given velocity is proportional to the amount of molecular gas in unit velocity width, N(Hl )/ d1,,and the OH opacity T()H (not T ~ , , , , ,but the real opacity) at a given velocity is also proportional to N(H2)/&:: thus T ~ ) H= ZT(.o where 2 is a constant. (3) The excitation temperature of OH, T,,,(OH), is uniform. (4) The 18 ciii continuum emission is optically thin and arises from a distributed, axisymmetric volume emissivity, j ( r ) [ r is the Galactocentric radius]. Now 2 and T,,(OH) are unknown parameters: how to determine them is described in $ 2.3. When a cloud whose OH opacity is T ( ) H is located at s = sg ( s is the position along the line of sight), the cloud absorbs the continuum intensity behind it, J"; j ( r ) d s . The absorption depth is written as ,f[1 ~ ~ ) ( - T ~ H ,j('r)d,s ) ] [ J- T,,(OH)] " : L where f is the beam filling factor of OH absorbing gas. The : j ( ~ ) d sand , the apparent opacity of the cloud is expressed as observed continuum intensity is:J -
and Tc.0(and thus 7()H = ZT&) are known through the observations, we can obtain the value Since 7EL,,p j ( r ) d s if we know f , 2,and Tpx(OH).Then s o is derived from the value of::J , j ( r ) d s using a -w distribution model of j ( r ) . We assume that the 18 cm continuum cmissivity , j ( r ) around the Galactic center is described as a sum of several axisymmetric Gaussians:
"fJ'""
Figure 3 shows the schematic relation of geometrical parameters. The Galactic longitude !is in degrees. Here T , s, and (T! are in units of projected distance corresponding to I" at the distance to the Galactic center: i.e., 150 pc at 8.5 kpc. The observed longitudinal distribution of the continuum brightness Tcont(P) can be written as Tcoll+(P) = ;:/, j ( ~ds. ) Since the continuum emissivity beyond the solar circle should j ( r )ds can be replaced with JTxj ( T ) d s . Thus be negligible compared with that in the center. ;:"J
Now n , and ~7~can be drawn so that the Eq. (3) reproduces the observed longitudinal distribution o f t h e continuum brightness. It is found that ohserved continuum distribution is well fitted by up to three components (iV = 3 ) . For b = 000,we obtain nl = 98.3, a 2 = 34.0, o y = 11.4; (TI = 0.120, g 2 = 0.677, (73 = 7.18 ( a zare in Kelvins, 0,are in degrees). Figure 4 shows the result of our fit. The model reproduces the observed data quite well. Since C T ~ 0, 2 are small enough (6 100 pc) and al,a2 are large, the continuum distribution due to the first ( i = 1) and the second ( i = 2) components is well defined. On the other hand, because of large ( ~ and ~ 1 small 0 3 , the third (i = 3 ) component should suffer from possible non-axisymmetric distribution on the largest scale andor individual sources; furthermore, small dT,.,,+/ds causes large errors in the derived positions of clouds. Therefore, i n longitudes where the contribution from the first and the second components is negligiblc (i.e., (el 2 105, corresponding to f 2 2 0 pc), the positions of clouds obtained using this model are rather uncertain.
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Line of Sight A
I’ lSO
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:loot h“ so
0
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-2
-4
-6
Galactic Longitude [degrees] Fig. 3 The schematic relation of gcoinetrical parameters in the face-on view of the Galactic center region.
Fig. 4 Longitudinal distribution of the 18 cm continuum in units of antenna temperature at b = O?O by Boyce & Cohen (1994) (open circles) and a fit by 3 Gaussians (solid line).
2.3 Choice of the 2 and T,,(OH) Values Following the procedure, we can draw a face-on (2-y) distribution of the molecular gas by putting each data point of the !-v diagram onto the x:-y plane with a projection of (P. s) + (:r. y). However, f, Z,and T,,(OH) are still unknown. Here we assumc ,f = 1. There is no bottom-up scheme to determine 2 and T,,(OH). We employ a trial-and-error scheme, making face-on maps at b = 0?0. 1t012. and *0?4 with various values of 2 and T,,(OH). Trials have been done for 2 = 0.04 to 0.70 [K-’1 and T,.,(OH) = 0 to 10 [K]. We have chosen an appropriate set OF (Z. T,-,(OH)) so that the following three conditions are satisfied: ( I ) The resultant face-on distribution of the CO brightness is not too asymmetric between the near and far sides with respect to the centcr; (2) The features extending above and below the Galactic plane are placed in similar face-on positions a1 different latitudes; and (3) Most of the CO emission has a solution for the position so. By combining these conditions, we have chosen Z = 0.15 & 0.03 [K-’] and T,,(OH) = 4 f l [K]. Sets of(larger 2,smaller T,,(OH)) and (smaller Z,largcrT,,(OH)) are rejected from conditions (1) and (3), respectively. If we adopt smaller or larger Trx(OH) values, condition (2) is not satisfied. The validity of the parameters, f,2 , and T,,(OH), is discussed in 5 3.3.
3 Results and Discussion 3.1 Overall Structure Figure S shows the resultant face-on view of the central 1kpc x 1kpc: of the Milky Way at b = 0?0, seen from the north Galactic pole. The CO brightness of each pixel in Fig. 2 is projected onto the y y plane and smoothed by m = 20 pc Gaussian. Figures Sn and Sh show the distributions of CO brightness and corresponding radial velocity, respectively. As seen in Fig. 2, high O H / C O ratios are more widespread at positive longitudes. This comes about because the face-on CO distribution is tilted so that the gas in the Galactic-eastern (positive 2 ) side lies closer to us, as qualitatively suggested by Cohcn & Few (1 976). An elongated molecular cloud condensation (“central condensation”), whose size is approximately 500 x 200 pc, dominates the CO emission in the Galactic center region. It should be compared with “twin peaks” in central regions of barred galaxies (Kenney et al. 1992). The minor axis length of the condensation might be smaller since the face-on map involves positional errors along the lines of sight. The major axis of the condensation is tilted with respect to thc line of sight (n. = 0) by E 70” so that the Galactic-eastern side
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Fig. 5 The resultant molecular face-on view 01 the Galactic center at b = 0’0 ( a ) Distnbution of CO J = 1 - 0 emission (normalized by the peak value). Contours are 0 01,0.02, 0 04, 0 08, 0.16, 0.32, and 0 64 of the peak value (h) Correvponding radial velocity The solar system is located dt (z, y) = (0. -8500)
is closer to us: this angle does not change significantly in the acceptable parameter space of (2. Tc:,(OH)). The condensation includes the Sgr A ( B E 0?0),Sgr B (B N 0?6), and Sgr C (! ru -0”) molecular cloud complexes and a huge molecular cloud complex at P r-. 1:s (“1:s region”). The face-on distribution of radial velocity (Fig. 5b) shows that the gas motion in the condensation is strongly noncircular. The gas on the far-side has larger receding velocity, while the gas on the near-side is approaching. This velocity field can be explained if the gas orbits are elongated along the major axis of the condensation. A similar trend is often seen in both numerical simulations of gas kinematics in a barred potential and observations of barred galaxies (see, e.g., Athanassoula 1992; Lindblad, Lindblad, & Athanassoula 1996). This fact agrees with the arguments that the Milky Way is barred. The EMR feature corresponds to the peripheral region surrounding the central condensation.
3.2
Nature of the Gas in the Central Hundreds of Parsecs
Oka et al. (1996) argued that formation of massive stars may have been taking place in the central 100 pc of radius (“star-forming ring”), based upon comparisons between their CO data, hydrogen recombination line emission (Pauls & Mezger 1975), and OWIR stars (Lindqvist et al. 1991). The “star-forming ring” forms a part of the central condensation. Sofue (1995) identified a pair of arm-like features in the central condensation from the I 3 C 0 data taken by Bally et al. (1987), though these “arms” are not clearly separated in our face-on map. This may he due to insufficient spatial resolution of the present analysis and to deviation from the assumed smooth, axisymmetric continuum emissivity because of embedded discrete sources. Considering Sofue’s two-arm model and our face-on map, active star-forming regions Sgr B and C are both located on the leads of the arms. Conccntration of clouds onto the arms would he occurring due to some kind of gas orbit crowding. The locations of the sites of active star formation on the leading edges of the inner arms may suggest a time lag between the arrival of the gas into the central orbit and the beginning of star formation as discussed by Kohno et al. (1999) for NGC 695 1. Our results are schematically summarized in Figure 6. We have found that the major axis of the central condensation is tilted by an angle of N 70” with respect to the line of sight. Therefore, the inclination between the major axes of the large-scale stellar bar and the central condensation is ru 40”-50”, if our viewing angle of the bar is ru 20”-30” (see Gerhard 1999, for reviews). In some barred galaxies, there are central molecular gas concentrations (“twin peaks”, Kenney et al. 1992). In a face-on projection of high resolution images, some of these concentrations tend to consist of N
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a pair of arms and their major axes are tilted with respect to the bar major axes by moderate angles: such as IC342 by lshizuki et al. (1 990), M 10 1 etc. by Kenney et al. ( 1 992) NGC 695 1 by Kohno et al. (l999), and NGC 4303 by Koda et al. (in preparation). They are similar to our face-on view of the Galactic center. Hydrodynamical simulations of the barred galaxies also produce similar gas distributions. Central Condensation /, .. A pair of arms7
c
/
SgrB
Observer
300 pc
Fig. 6 A picture of molecular gas in the Galactic center region proposed on the basis ofthe results from this work. Linear scales are approximate.
3.3
I [pc]
Fig. 7 Contours show the integrated continuum emissivity j ( i - ) d s at b = 0eO.Levels are 2.5, 5.0, 7.5, 10.0, and 20.0 K in units of brightness temperature. The center (I-,y) = (0,O) is shown by a cross. Dashed lincs indicate I = +1?5.
Validity of Parameters
Beam filling factor. We have assumed that f , thc beam filling factor of the OH absorbing gas, is unity. From Eq. ( l ) , T ~ ~gives , ~ , the lower limit o f f . For some lines of sight, T:~,,,, is rather large: 0.64 toward Sgr B; 0.35 toward the 1:s region and Bania's Clump 2 (Bania 1977) at B r , 3". Sawada et al. (2001) estimated the beam tilling factor of CO J = 1 - 0 emission to be 0.4-0.7 based upon a large velocity gradient analysis and high resolution CO data taken by Oka et al. (1998b). Since it is expected that the OH absorption arises also from less densc cloud envelopes compared with thc CO emission, thc beam filling factor of OH absorbing gas is at least similar and can be even largcr. These facts indicate that f is large and can be reasonably replaced with unity. The "Z" factor. There are three previously-known relations. ( I ) The column density of molecular hydrogen N(H2) is proportional to the CO .J = I - 0 integrated intensity. The conversion factor (X-factor) is measured to be about 2 x 102"[ ~ i i i (K - ~kiris-')-'] for molecular clouds in the Galactic disk (see, e.g., Dame, Hartmann, h Thaddeus 2001). For the Galactic center clouds, however, it is reported that the X-factor is smaller than that for the Galactic disk: i.e., X = scveral x 10'' (see, e.g., Oka et al. 1998a). (2) The relative abundance of OH to molecular hydrogen, [OH]/[H2], is typically sr,vpra,l x lop7 (see, e.g., Herbst & Leung 1989). (3) The OH column density is derived from the OH 1667MHz opacity 7 as N(OH) = 2 x 1014T,, / T ~ U Using . these relations, the 2 factor is written as Z [K-l] = 5 x lo-:' ([OH]/[H2]/10-7) ( X [cnir2( K k ~ i i s - ' ) - ~ ] / l O (lO/T,, ~~) [K]). This value becomes several x lo-', several times smaller than our value, 0.15. This discrepancy can be understood if we consider the higher OH abundance (sevcml x lo-"), which may be due to interstellar shocks (see, e.g., Wardle 1999) and/or a high ionization rate caused by X-rays (Lepp h Dalgarno 1996). Excitation temperature. Figure 7 shows the integrated continuum emissivity, j ( r ) d s in Eq. ( I ), at b = 0?0. It is expected that at each location in the iigure, any gas whose T,,(OH) is higher than the
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integrated continuum emissivity at that location would be seen in emission. However, no significant OH emission is observed. Thus, if T,,(OH) 2 (i[K],the face-on gas distribution is almost restricted to near side of the Galactic center, g < 0, which is unrealistic. On the other hand, Boyce & Cohen (1994) noted that the OH absorption was only detected at positions where the continuum temperature exceeded 5 K of antenna temperature (7.5 K of brightness temperature). Lower limit of T,, (OH) is provided from this fact: T,,(OH) is higher than the integrated continuum emissivity, which is typically half of the observed continuum brightness for the Galactic center clouds. Thus T,,,(OH) should not be far below (7.5/2) K. These two constraints are consistent with our value, 4 K.
4
Conclusions
We have developed a method to derive positions of molecular clouds along the lines of sight. The method is completely independent of any kinematic model and based on observable data alone; the C O emission line, the OH absorption line, and 18 (:m continuum distribution. It is applied to the central region of the Milky Way to obtain a molecular face-on map. Most of the CO emission comes from the “central condensation”. It is elongated, and its major axis is tilted wilh respect to the line of sight by = 70” so that the Galacticeastern end is closer to us. The gas within it shows highly noncircular motion: the gas in the far side is receding whereas the gas in the near side is approaching. This noncircularity of the gas motion is most likely induced by a barred potential. The results give a new evidence for the existence of a bar in the Milky Way Galaxy based on direct distance derivation independent of kinematic models. Acknowledgements We acknowledge Leonardo Bronfman for providing us the CO d = 1 - 0 data in a computer readable form This work was supported by a Grant-in-Aid for Scientific Research of the Ministry of‘ Education, Culture, Sports, Science, and Technology 08404009 and 10147202.
References Athanassoula, E. 1992, MNRAS, 259, 345 Bally, J., Stark, A. A., Wilson, R. W., & Henkcl, C. 1987, ApJS, 65, 13 Bania, T. M. 1977, ApJ, 216, 381 Binney, J., Gerhard, 0. E., Stark, A. A., Bally, J., & Uchida, K. I. 1991, MNRAS, 252, 210 Bitran, M., Alvarez, H., Bronfman, L., May, J., & Thaddeus, P. 1997, A&AS, 125, 99 Boyce, P.,& Cohen, R. J. 1994, A&AS, 107, 563 Cohen, R. J., & Dent, W. R. F. 1983, in Surveys of the Southern Galaxy, ed. W. B. Burton & F. P. Israel (Dordrecht: Reidel), 159 Cohen, R. J., & Few, R. W. 1976, MNRAS, 176,495 Dame, T. M., Hartmann, D., & Thaddeus, P. 2001, ApJ. 547,792 Gerhard, 0. E. 1999, in ASP Conf. Ser. 182, Galaxy Dynamics, ed. D. Memtt, M. Valluri, & J. Sellwood (San Francisco: ASP), 307 Herbst, E., & Leung, C. M. 1989, ApJS, 69, 27 I Ishizuki. S., Kawahe, R., Ishiguro, M., Okumura, S . K., Morita, K.-I., Chikada, Y., & Kasuga, T. 1990, Nature, 344,224 Kaifu, N., Kato, T., & Iguchi, T. 1972, Nature Phys. Sci., 238, 105 Kenney, J. D. P., Wilson, C. D., Scoville, N. 2.. Devereux, N. A,, & Young, J. S. 1992, ApJ, 395, L79 Kohno, K., Kawabe, R., & Vila-Vilar6, B. 1999, ApJ, 5 I I , 157 Lepp, S., & Dalgamo, A. 1996, A&A, 306, L21 Lindblad, P. A. B., Lindhlad, P. O., & Athanassoula, E. 1996, A&A, 313, 65 Lindqvist, M., Winnherg, A., Hahing, H. J.. & Matthews. H. E. 1991, A&AS, 92, 43 Liszt, H., & Burton, W. B. 1980, ApJ, 236, 779 Oka, T., Hasegawa, T., Handa, T., Hayashi, M., & Sakamoto, S. 1996, ApJ, 460, 334 Oka, T., Hasegawa, T., Hayashi, M., Handa, T., & Sakamoto, S. 1998a, ApJ, 493, 730 Oka, T., Hasegawa, T., Sato, F., Tsuboi, M., & Miyazaki, A. 1998t1,ApJS, I 18, 455 Pauls, T., & Mezger, P. G. 1975, A&A, 44. 255, Sawada, T., et al. 2001, ApJS, 136, 189 Scoville. N. Z. 1972, ApJ, 175, L127 Sofue, Y. 1995, PASJ, 47, 527 Wardle, M. 1999, ApJ, 525, LlOl
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Astron. Nachr./AN 324, No. SI, 93-99 (2003) I DO1 10.1002/asna.200385109
The Inner 200pc: Hot Dense Gas Christopher L. Martin* I, Wilfred M. Walsh', Kecheng Xiao', Adair P. Lane', Christopher K. Walker', and Antony A. Stark'
' Harvard-Smithsonian Center for Astrophysics, 60 Garden St., MS-12, Cambridge, MA 02138 ' Steward Observatory,University of Arizona, 933 N. Cherry Ave., Tucson, AZ 85721 Key words Ga1axy:center 1SM:molecules
-
Ga1axy:kinematics and dynamics - 1SM:atoms - 1SM:general
-
Abstract. We present fully-sampled maps of 461 GHz CO J = 4 --7' 3 , 807 GHz CO J = 7 + 6, and 492 GHz [C I] 'PI --t "POemission from the inner 3 degrees of the Galactic Center region taken with the Antarctic Submillimeter Telescope and Remote Observatory (AST/RO) in 2001-2002. The data cover -1P3 < e < 2', -0P3 < b < 0y2 with 0.5' spacing, resulting in spectra in 3 transitions at over 24,000 positions on the sky. The CO .I = I 3 emission is found to be essentially coextensive with lower-J transitions of CO. The CO J = 7 i 6 emission is spatially confined to a far smaller region than the lower-J CO lines. The [C 11 'PI + 'PO emission has a spatial extent similar to the low-J CO emission, but is more diffuse. Bright CO d = 7 + 6 emission is detected in the well-known Galactic Center clouds SgrA and SgrB. Analyzing our CO J = 7 + 6 and CO J = 4 + 3 data in conjunction with J = 1 + 0 " C O and I3 C/O data previously observed with the Bell Laboratories 7-m antenna, we apply a Large Velocity Gradient (LVG) model to estimate the kinetic temperature and density of molecular gas in the inner 200 yc of the Galactic Center region. Typical pressurcs in the Galactic Center gas are n(EI2) . Tklr, 10" K cin-". We present an % !( b ) map of molecular hydrogen column density derived from our LVG results.
-
7
-
1 Introduction
-
Much has been learned about dense gas in the Galactic Center region through radio spectroscopy. Early 2) OH absorption (Robinson et al.(l964),Goldstein et aL(1964)) suggested the observations of F ( 2 existence of copious molecular material within 500 pc of the Galactic Center. This was confirmed by detection of extensive J = 1 4 0 "CO emission (Bania(lY77),Liszt & Burton(1978)). Subsequent CO surveys ( Bitran(1987), Stark et a1.(1988), Oka et a1.(1998), Bitran et a1.(1997)) have measured this emission with improving coverage and resolution-these surveys show a complex distribution of emission, which is chaotic, asymmetric, and non-planar; there are hundreds of clouds, shells, arcs, rings, and filaments. On scales of 100 pc to 4 kpc, however, the gas is loosely organized around closed orbits in the rotating potential of the underlying stellar bar (Binney et al.( 1991)). Some CO-emitting gas is bound into clouds and cloud complexes, and some is sheared by tidal forces into a molecular inter-cloud medium of a kind not seen elsewhere in the Galaxy (Stark et al.( 1989)). This diffuse inter-cloud medium appears in absorption in F ( 2 + 2) OH (McGee(1970),Robinson & McGee(l970)), in (110 + 111) H2CO (Scoville et al.(l972)), and in J = 0 + 1 HCO+ and HCN (Linke et aL(l981)). In contrast, the clouds and cloud complexes are dense, as they must be to survive in the galactic tide, and they appear in spectral lines which are tracers of high density (n(H2) > lo4 (m-')), such as NH3 (1, 1 ) (Giisten el aL(1981)) and CS cJ = 2 + 1 (Bally et al.( 1988)). The large cloud complexes, Sgr A, Sgr B, and Sgr C, are the among the largest molecular cloud complexes in the Galaxy ( A 1 2 10"' M a ) . Such massive clouds must be sinking * Corresponding author: e-mail:
cmartin @cfa.harvard.edu @ 2003 WILtY-VCH Vrrlag GwhH B Co K G . A Weinhem
C. Martin et al.: AST/RO: Inner 200pc
94
toward the center of the galactic gravitational well as a result of dynamical friction and hydrodynamic effects (Stark et al.( 1991)). The deposition of these massive lumps of gas upon the center could fuel a starburst or an eruption of the central black hole (Genzel & Townes( 1987)). As prelude to further study of the Galactic Center molecular gas, we would like to determine its physical state-its temperature and density. This involves understanding radiative transfer in CO, the primary tracer of molecular gas. Also useful is an understanding of the atomic carbon lines, [C I], since those lines trace the more diffuse molecular regions, where CO is destroyed by UV radiation but Hz is still present. The .J = 1 + 0 I 2 C 0 line is often optically thick. Its optical depth can be estimated by studying its isotopomers, '"CO and C " 0 . In the Galactic Center region, l'C0 is 24x less abundant than l 2 C 0 (Penzias( 1980),Wilson & Matteucci(l992)), and C l 8 0 is 250x less abundant than 12C0 (Penzias(1981)). Since the radiative and collisional constants of all the isotopomers are similar, the ratio of optical depths in their various spectral lines should simply reflect their relative abundances. Where the lines are optically thin, the line brightnesses should be in the same ratio as the isotopic abundances; where the lines are thick, deviations from the abundance ratios are a measure of optical depth. Bally et a].(19871, Bally et al.( 1988) and Stark et al.( 1988) produced fully-sampled surveys of "CCO and '"C0 in the Galactic Center region. They find that the ratio of the " C O .I = 1 + 0 to "CO J = 1 + 0 line brightness temperatures (TtZo/T&) is typically 10 3:2 in Galactic Center gas that is far from dense cloud cores. This a),especially in indicates much of the Galactic Center l 2 C 0 emission is only moderately thick 6 (Polk et al.( 1988)) is smaller, comparison to the galactic disk outside 3 kpc radius, where T;:"/T;:(, even though the isotope ratio 12C/13C N 40 (Penzias( 1980)) is larger. Heiligman( 1982) and Dahmen et aL(1998) made surveys in ClSO .J = 1 40. These show "C0 J = 1 -+ 0 to C"0 J = 1 0 line brightness temperature ratios (T:?+o/T;50)which vary from 40 to over 200, with typical values near 70, indicating values of T::(, which vary from 3 to less than I , while the core region of SgrB2 shows 7;:" 10. Determining the excitation temperature of CO works best if emission lines from several J levels have been measured. Lacking such observations, what is often done is to use the brightness temperature of the "CO J = 1 40 line as a lower limit to the excitation temperature of the J = 1 state, T e x , ~ =This l. estimate can be misleading, because the emission may not fill the telescope beam, diluting the brightncss temperature and causing it to be many times smaller than Tex,.~=l; as will be apparent from the data to be presented here, this is thc usual case for gas in the Galactic Center region. Moving up the energy ladder, Sawada et a1.(2001) surveyed the Galactic Center region in l 2 C 0 J = 2 4 1. They compare their data to the .J = 1 0 data of Bitran et aL(1997) and find Ti51/T:?+o = 0.96 0.01, with little spatial variation. What this means is that almost all the CO in the Galactic Center region has l o w 4 states which are close to local thermodynamic equilibrium (LTE), so that the excitation temperatures Tex, J of those states are all close to the kinetic temperature, Tki,, and , the ratio of line brightnesses for transitions between those states are near unity and therefore independent of Tkill (cf. Goldreich & Kwan(l974)). LTE in the low-.J states of CO does not occur under all circurnstanccs in the interstellar medium, but it is very common and appears to be the rule for Galactic Center gas. For each value of Tkin and ~ ( H Z )there , will, however, be some value of J above which all higher-J states fail to be populated, because their Einstein A coefficients (which increase as J3)are so large that the collision rate at that value of Tki,, and n(H2) cannot maintain those states in LTE, and they must therefore be subthermally excited, i.e., T e x ,)
( 51
where LCOi s in K km spl pc2, and the uncertainties are 1c. The exponent of the LCO-MVTrelation is the similar to that of the disk clouds (= 0.81).
T. Oka and T. Hasegawa: Gravitational Stability of Molecular Clouds
1 04
I o9
100
I ox 10' c
0
lo6
2
3
10'
I o4
0. I 0.1
10
1
s
100
LCO [K km s-1 pcz]
[PCI
Fig. 2 A S-uv plot is shown for the clouds identified both in the 45 m (filled polygons) and 1.2 m (open circles) data sets, with the same plots for the disk clouds (dots) and the large GC clouds identified manually in the CO J=2-1 survey data (crosses). Shapes of filled polygons indicate boundary intensities, T,,,,, = 5 K (filled circles), 7.5 K (filled rhombus), and 10 K (filled triangles).
4
Fig. 3 A LCO-MVTplot is shown for the clouds identified both in the 45 m (filled polygons) and 1.2 rn (open circles)data sets, with the same plots for the disk clouds (dots) and the manually identified large GC clouds (crosses). Shapes of filled polygons indicate boundary intensities, T,,,, = 5 K (filled circles), 7.5 K (filled rhomhus), and 10 K (filled triangles).
Discussion
4.1 Cloud in Equilibrium with External Pressure The larger MVTvalues of the GC clouds than those of the disk clouds with the same LCOcan not be solved by introducing a large CO-to-Hz conversion factor, since y-ray (Blitz et al. 1985), far-infrared (Sodroski et al. 1995), and X-ray observations (Sakano et al. 1997) suggest CO-to-HZ conversion factors in the Galactic center smaller than the standard value. ~ can be understood in terms of model clouds in pressure The large variations in the L ~ o - b f "plot equilibrium with an intercloud medium (Oka et al. 1998a). The virial equation for a spherically symmetric non-magnetic cloud of mass M and radius R, embedded within an intercloud medium of pressure p can be written as 1 dzI -= 3a;M 2 dt2
~
G M ~ - 4nR"p R
a-
where I is the generalized moment of inertia of the cloud, ov is the velocity dispersion of turbulent motion (thermal motion has been neglected), and a is a dimensionless coefficient of order of unity. The virial equation can easily be extended to the case of a magnetized cloud by substituting a to a,if = a[l (@/@cr)2(1 - R/Ro)](Nakano 1998). In the equilibrium state, d 2 1 / d t 2=O, we have two equilibrium masses ,
where j 3 = 16.rraR2p/9a$ is a dimensionless parameter, and Mo = YRrr$/G is the commonly used virial theorem mass which is the same as Afv, in eq.(3). /? is also expressed by three dynamical parameters as
105
Astron. Nachr./AN 324,No. S 1 (2003)
( P / T ) ,where U is the gravitational potential energy, T is the internal kinetic energy, and P = 47rR3p . It is easily confirmed that when the gas behaves isothermally, an equilibrium state with a+
[ j = ( / U I / T )x
(> 1/2) is gravitationally unstable and that with
(1- (< l/Z) is gravitationally stable. This N denotes the degree of gravitational instability. For an opaque cloud with uniform CO brightness of TCO, the total CO luminosity can be expressed as Lpo = 27rR2Tcoo1,,. From this expression, and equating the virial mass &fvT = trl& to (4/3)7rR3pp, where p is the mean mass density, we get the linear Lco-A'lv~ relation,
Introducing the CO-to-Hz conversion factor X = N(Hz)/Ico, on the other hand, the total molecular mass including the helium correction 1.36 is written as
M
= 2.2 x 1 0 - 2 " x L c o
(&J)
(10)
where Lc0 is in K km pc2. Using the virial mass Afv, = a h f , as the total molecular mass, and comparing eqs.( I 1) and (12), we get
x
=
a+xo,
X o stays constant when rewritten as Mo
= 2.2 x
r
)
p x / T c ~ doe5 not vary in the mean from cloud to cloud. Thus eq.( 1 1 ) is
lO-*"tr-fXoLco.
The large variation in the LCO-MVT relation for each region or scale. 4.2
(1 3 ) in
eq.( 15) is explained by choosing different
(Y
in the eq.( 15)
Gravitational Stability
A theoretical study has shown that interstellar clouds which are close to gravitational instability exhibit precisely the same scaling laws, provided that they are in equilibrium with a constant pressure environment (Chikze 1987). The disk clouds actually follow the same scaling laws, defining a base of highly scattered distributions (Figs.2, 3). The paucity of clouds below the Lco-MvT trend of the disk clouds suggests that such clouds are gravitationally unstable ((Y > 1/2) and may have moved to stable zone by contraction or have collapsed to form new stellar clusters. Assuming that the disk clouds are nearly at the onset of (a linear relation between R a n d gravitational instability, ( Y = 1/2, and using AJxJ, defined by eq.(3) as M,, cm-2 (K km s-')~'. This gives a the size parameter, S, is implicitly assumed), we get X O= 1.6 x CO-to-Hz conversion factor X = 1.2 x 10'" cmP2 (K km s-l)-' for the disk clouds, which is smaller than the standard value by a factor of 3, but is close to that derived from the y-ray observations of the Orion molecularcloud, X=1.06 x em-' (K km s - ' ) ~ ' (Digel, Hunter, & Mukherjee 1995). Although we do not have compelling evidence that the disk clouds are nearly at the onset of gravitational instability, diagnosis of gravitational stability by the parameter (1 must be meaningful in a reiative sense. We diagnose the degrees of gravitational stability of the CC clouds by (Y using X" = 1.6 x 102" cm-2 ( K km s - l ) - ' . All but one G C clouds have (1 smaller than 1/2, which suggests they are gravitationally stable. They might be in equilibrium with the external pressure of intercloud medium. This fact may explain the paucity of star formation activity in this region. We show the 1-V distributions of the GC clouds with three a ranges separately (Fig.4). Clouds with (Y > 0.01 follow the main ridge of intense CO emission, part of which defines two rigidly-rotating molecular
106
T. Oka and T. Hdsegawa: Grrlvitdtional Stability of Molecular Clouds
arms (Sofue 1995). These molecular arms are associated with a ring of HI1 regions, which is called the star forming ring (Oka et al. 1996). On the other hand, clouds with cy < 0.01 spread over the 1-V plane irrespective of the location of the Sofue’s molecular arms. quit
-v I)
E
200 t
-I
100
Y
0 -100 -200
-*.*
2.0
1.5
1.0
0.5
0.0
4.5
Galactic Longitude [degrees]
j -1.0
.
0
1.5
1.0
0.5
0.0
j
2‘ ”
-0.5
’
-1.0
Galactic Longitude [degrees]
Fig. 4 The l-V distribution of the 45 m GC clouds with ( a ) a 2 0.01 and (b)a < 0.01. Shaded areas denote Sofue’s molecular arm I-IV (Sofue 1995).Broken lines show the 1-Vloci of the Galactic a r m .
Figure 5 shows the correlation plots between basic cloud parameters and their degrees of gravitational = 5 K sample of the 45 m GC clouds. The largest five clouds with S 2 20 pc, instability for the Trnin which might be categorized as cloud complexes, were indicated by open circles. No correlation is seen in the S - n or Lc0-a plots, while gv correlates inversely with except for the largest five clouds. A anti-correlation in h f ” ~ - plot n is a result of the r v - u anti-correlation. These suggest that the degree of gravitational instability of the GC clouds is intimately related to the velocity dispersion. Neither size nor mass is the determining factor for gravitational instability. These results imply that gravitational instability of the GC clouds grows with dissipation of velocity dispersion. We conclud that dissipation processes associated with the star forming ring promote gravitational instability of clouds and subsequent star formation. Mechanisms such as orbit crowding at the inner Lindblad resonance may play a role. lkuta & Sofue (1997) studied kinematic properties of giant molecular clouds (GMCs) associated with star forming regions in the Galaxy, and found that the star formation efficiency ( S F E ) is inversely correlated with the velocity dispersion and the virial mass of a GMC. Kohno et aL(1999) found, in the central region of the barred spiral galaxy NGC 6951, an anti-correlation between the dense molecular gas fraction and the gas velocity dispersion. They claimed the formation of dense molecular gas is caused by gravitational instability, which is intimately related to the dissipation of velocity dispersion. These results also imply that the dissipation of random motion within a GMC promotes the gravitational instability and subsequent star formation.
5
Conclusions
Our analyses of the large-scale CO data by the NRO 45 m telescope have yielded samples of molecular clouds in the Galactic center region. The major results of our analyses are the following: 1 . All clouds and cloud complexes in the Galactic center, exept for one cloud, are gravitationally stable, being in equilibrium with the external pressure of intercloud medium. 2. Less stable clouds in the Galactic center concentrate on the Galactic plane, following the main ridge of intense CO emission, part of which defines the two rigidly-rotating molecular arms.
Astron. Nachr./AN 324, No. S 1 (2003)
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Fig. 5 Relations between size, velocity dispersion, CO luminosity, virial mass, and the parameter u1 for the T,,,,, = 5 K sample of the 45 m GC clouds. Open circles indicate the largest 5 clouds with S > 20 pc, which might be categorized as cloud complexes.
3. The velocity dispersion of a cloud correlates inversely with the degree of gravitational instability. Dissipation processes associated with the star forming ring may promote the gravitational instability of molecular clouds and subsequent star formation. telescope key program,
References Blitz, L., Bloemen, J. B. G. M., Hermsen, W., & Bania, T. M. 1985, A&A, 143, 267 Chi&, J. P. 1987, A&A, 171,225 Digel, S. W., Hunter, S. D., & Mukherjee, R. 1995, ApJ, 441,270 Ikuta, C., & Sofue, Y. 1997, PASJ, 49, 323 Kohno, K., Kawabe, R., & Vila-Vilaro, B. 1999, ApJ, 51 1, 157 Miyazaki, A,, & Tsuboi, M. 2000, ApJ, 536,357 Nakano, T. 1998, ApJ, 494,587 Oka, T., Hasegawa, T., Hayashi, M., Handa, T., & Sakamoto, S. 1996, ApJ, 460, 334 Oka, T., Hasegawa, T., Hayashi, M., Handa, T.. & Sakamoto, S. 1998a. ApJ, 493,730 Oka, T., Hasegawa, T., Sato, F., Tsuboi, M., & Miyazaki, A. 1998b, ApJS, 118, 455 Oka, T., Hasegawa, T., Sato, F., Tsuboi, M., & Miyazaki, A. 2001a, PASJ, 53, 779 Oka, T., Hasegawa, T., Sato, F., Tsuboi, M., & Miyazaki, A. 2001b, PASJ, 53, 787 Oka, T., Hasegawa, T., Sato, F., Tsuboi, M., Miyazaki, A., & M. Sugimoto 2001c, ApJ, 562, 348 Sakano, M., et al. 1997, IAU Symp. 184, The Central Regions of the Galaxy and Galaxies, 227 Scoville, N. Z., Yun, M. S., Clemens, D. P., Sanders, D. B., & Waller, W. H. 1987, ApJS, 63, 821 Sodroski, T. J., et al. 1995, ApJ, 452,262 Sofue, Y. 1995, PASJ, 47,527 Solomon, P. M., Rivolo, A. R., Barret, J., & Yahil, A . 1987, ApJ, 319, 730 Young, J. S., & Scoville, N. Z. 1991, ARA&A, 29, 581
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Astron. Nachr./AN 324. No. SI. 109- 115 (2003)/ DO1 10.1002/asna.200385056
Spectroscopy of Hydrocarbon Grains toward the Galactic Center and Quintuplet Cluster J.E. Chiar*'.*,A.J. Adamson3,D.C.B. Whittet4,and Y.J. Pendleton' I
NASA Ames Research Center, Mail Stop 245-3. Moffett Field, CA 94035
* SET1 Institute, Mountain View, CA 94043
Joint Astronomy Centre, 660 N. A'ohoku Place, University Park, Hilo, Hawaii 96720 Rensselaer Polytechnic Institute, Department of Physics, Applied Physics, and Astronomy, Troy, NY 12180
Key words extinction, interstellar medium dust and molecules, circumstellar dust Abstract. Our view of the Galactic center (GC) is affected by extinction from both diffuse interstellar medium (ISM) dust and dense molecular clouds along the line of sight. The enormous visual extinction present toward the center of our Galaxy (- 31 magnitudes) necessitates a study of the interstellar dust properties as well as an investigation into the distribution of the different dust components. We have built upon the historic spectroscopy o f Willner et al. (19791, Butchart et al. (1986), and McFadzean et al. (1989) in order to investigate the distribution of these dust components across the GC field. Specifically, we employ spectroscopy in the 3 pm region to investigate absorption features at 3.0 pm and 3.4 pm in lines of sight toward the GC central cluster and the Quintuplet cluster to the northeast. The 3.4 pm feature is one of the primary spectral signatures of the organic component of interstellar dust and is, to date, only observed in the cold diffuse interstellar medium. The 3.0 ice feature is carried by dense molecular cloud material, and can therefore be used to loosely trace the distribution of such material across the GC field. By obtaining spectra for multiple sightlines we have been able to deconvolve the diffuse ISM and dense molecular cloud components. Our study shows that differences exist in the spectra of relatively nearby lines of sight in the Galactic center central cluster. The depth o l the 3.0 pm water-ice feature varies by a factor of almost 5 across a 2 parsec (in projection) region, perhaps reflecting the clumpy nature of the dense clouds. In addition, we found that the 3.4 pm hydrocarbon feature varics in depth across the areas studied toward the central cluster, whereas the depth is relatively constant toward the Quintuplet cluster. This is likely a reflection of the distribution of extinction from the foreground diffuse ISM. Our ground-based and space-based spectroscopy reveals differences in absorption features in the 3 and 6 pm regions between sightlines toward the GC central cluster and those toward the Quintuplet cluster. While the 3 pm spectra of both regions show a broad absorption feature blueward of the 3.4 pm absorption, only the Quintuplet spectra show a distinct absorption feature at 3.28 pm. This feature is indicative of the presence of polycyclic aromatic hydrocarbons (PAHs) along the line of sight. The Quintuplet-proper sources have 6 p m spectra that are markedly different than that of GC IRS 7 in the central cluster, and instead strongly resemble the spectra seen toward dusty late-type carbon-class (WC) Wolf-Rayet stars. This is the first hint of some spectroscopic similarity between the Quintuplet sources and dusty WC stars.
1 Introduction T h e center of our Galaxy is obscured by some 30 magnitudes of visual extinction. Investigations of objects located near the Galactic center (GC) require an understanding of the physical properties of the intervening dust through which all observations are made. Some of the dust is local to the G C , while much of it lies along the line of sight. Bright infrared sources located at the G C are used to probe this dust. T h e line of sight toward the G C is dominated by diffuse ISM dust (Lebofsky 1979), however it was suggested by * Correaponding author: e-mail: chiarQmisty.arc.nasa.gov, Phone: + I 6506040324, Fax: +1 6506046779 0xm WILEY VCH V C T I ~G~m m
co K G A wrlnhrlt~~
Chiar et al.: Spectroscopy of Dust Components toward the Galactic Center
110
McFadzean et al. (1989) that there is a molecular cloud component that contributes to the extinction toward the GC. Space-based spectroscopy with the Infrared Space Observatory (ISO) confirmed the presence of molecular cloud material by revealing absorption due to icy grain mantles (de Graauw et al. 1996; Lutz etal. 1996; Gerakines et al. 1999). It has been estimated that as much as one-third of the visual extinction arises in molecular cloud material (Whittet et al. 1997). Much of the molecular cloud extinction presumably arises in dust clouds located within 4 kpc of Earth (Sanders, Scoville & Solomon 1985). If this is so, then the clouds are not associated with the infrared sources that provide the continuum against which molecular absorption features can be observed. The diffuse interstellar medium is devoid of ices and instead contains only the refractory grain component, which includes aromatic [ring-like) and aliphatic (chain-like) hydrocarbons and silicates (Figs. 1 and 4). Hydrocarbons in the form of aromatics and aliphatics are a significant component of the diffuse ISM (e.g., Pendleton & Allamandola 2002). Aromatic hydrocarbons, observed throughout our Galaxy and other galaxies, are characterized by a family of bands normally observed in emission around 3.3, 6.2, 7.7, 8.6, 1I .3 and 12.7 pm (Allamandola, Tielens, & Barker 1989). Short-chained aliphatic hydrocarbons [Fig. 4, rightmost insert) are characterized by absorption at 3.4 pm and subfeatures due to CH2 (methylene) and CH3 (methyl) stretching modes at 3.38 and 3.48 pm (methylene) and 3.42 pm (methyl) (Sandford et al. 1991). Their spectral signature is seen not only toward the GC (Butchart etal. 1986), but in diffuse ISM dust throughout our Galaxy (Pendleton et al. 1994) and other galaxies (Wright et al. 1996). The 3.4 pm absorption feature observed in the diffuse ISM is distinct from that seen in dense molecular clouds (Brooke et al. 1996; Chiar et al. 1996). The latter is a smooth feature centered at 3.47 pm and the camer is thought to coexist with H2O-ice in the grain mantle. In contrast, the diffuse ISM hydrocarbons are most likely carried by a population of very small unaligned grains, rather than refractory mantles on silicate cores (Chiar et al. 1998; Adamson et al. 1999; Ishii et al. 2002), although the interpretation of these results remains controversial (Li & Greenberg 2002). I
I
I
I
I
I
! I Hydrucarbons
Dense Cloud Material
------
Diffwe Interstellar Medlurn
t
!
I 5
I 10
1
Wavelength ( p i )
Fig. 1 ISO-SWS spectrum centered on CC IRS 7. The beam size was 14” x 20“. Absorption features arising in the diffuse ISM and dense molecular cloud material are noted. Figure adapted from Lutz et al. 1996 ‘and Chiar et al. 2000.
2 Dense Cloud Absorption Features Ices in the cold molecular cloud material along the GC line of sight are characterized by absorption features due to HzO (3.0, 6.0 prn), COz (4.27 and 15.3 pm), and CH4 (7.67 pm) (Chiar etal. 2000; Gerakines et al. 1999; de Graauw etal. 1996; Lutz etal. 1996). The Short-Wavelength Spectrometer (SWS) on IS0
Astron. Nachr./AN 324, No. S1 (2003)
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h
0
In
0 ) --28
59 00
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17 42 34 17 42 32 17 42 30 17 42 28 17 42 26
Right Ascension (1950)
Fig. 2 Positions of infrared sources overlaid on a map o f HCN J= 1-0 emission (solid contours) and ionized gas at IS GHz (dottedcontours) [Reproduced from Giisten et al. 1987.1. HCN emission traces the high density molecular gas in the circumnuclcarring. Contour interval for the velocity-inlegratedHCN emission is 0.15 K averaged over 300kms-’ i n a 2” beam. Dotted contours of 15 GHz emission at 20 K intervals in a 3.6” x 3.4” beam. The radio-source SgrA* is located at cr = 17”42’”29.2” 6 = -28’59’19’’ The GC Quintuplet sources are located 14’ northeast of the GC, 30 pc away (assuming a distance of 8.5 kpc)
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observed the line of sight toward GC IRS 7 with a 14” x 20” beam. The spectrum from 2.6 to 12.5 pm is shown in Fig 1. Absorption features due to solid C 0 2 and CH4 are present; analysis of the profiles showed that the abundances of these molecules relative to HsO-ice are similar to those observed in local molecular clouds (Boogert et al. 1998; Gerakines et al. 1999; Chiar et al. 2000). The 6.0 y m absorption feature has been previously studied with the Kuiper Airborne Observatory (using FOGS and HIFOGS: Willner et al. 1979; Tielens et al. 1996), then later by ISO-SWS (Chiar et al. 2000; Fig. 5 ) . There is some controversy surrounding the precise identification of the 6 pm profile. although it is generally accepted that it can at least be partially attributed to HyO-ice with possible trace amounts of other ices (Chiar et al. 2000). A comparison between the 6 prn absorption profile observed toward G C IRS 7 and the Quintuplet-proper sources is shown in Fig. 5 and discussed in section 4. While airborne and space-based observations enabled us to study the average ice properties along the line of sight, spatial information could not be ascertained. The ISO-SWS observations, centered on GC IRS 7, were made with a large beam and included the sources GC IRS 1, IRS 3, much of the ionized bar and the ionized northern arm (Lutz et al. 1996). The KAO observations were carried out with a similarly large beam with FOGS and HIFOGS ccntered on GC IRS 3 (Tielens et al. 1996). In order to gain insight into the distribution of diffuse ISM dust and dense cloud material, we undertook a program of 3 Itm spectroscopy of the positions shown in Fig. 2 (Chiar et al. 2002) at the United Kingdom Infrared Telescope (UKIRT) using the the 0.6” slit of the Cooled Grating Spectrometer, CGS4. The 3 pm region includes absorption features due to H20-ice and aliphatic and aromatic hydrocarbons. Aliphatic hydrocarbons reside in the diffuse ISM and are discussed further below. Water-ice is the most abundant mantle constituent in any dense cloud, so we use it to trace the dense cloud material.
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Our spectroscopy shows that the HzO-ice profile shape is remarkably consistent across the seven lines of sight studied, including GC IRS 8 and IRS 19, which are located apart from the central cluster, closer to the GC circumnuclear ring (Fig. 2; Geballe et al. 1989). Compared to local molecular cloud ice profiles, e.g., in Taurus, the GC profile is broader and peaks at shorter wavelengths (Fig. 3). Laboratory spectra of pure HaO-ice are not able to account for the observed absorption in the GC ice profile; additional absorption is present shortward of 3 pm and from 3.2 to 3.6 p m (in addition to the 3.4 pm hydrocarbon feature). The possibility that the blue excess is due to NH3-ice has been discussed by Chiar et al. (2000), although other explanations may be plausible (Dartois & d'Hendecourt 2001). We find that the optical depth of the ice band is the greatest toward GC IRS 19 ( 7 3 . 0 = 1.5), a factor of almost 5 greater than the weakest ice band (Chiar et al. 2002). The variation of the 3.0 pm profile depth across the central GC field has been noted previously by McFadzean et al. (1989). We also obtained a 4.5-5 pm spectrum of the line of sight toward GC IRS 19, to search for evidence of solid CO absorption, since it has the deepest H2O-ice band (and therefore a high abundance of icy grain mantles). In addition to unresolved gas-phase CO lines, our spectrum shows a weak solid CO feature at 4.67 pm and an X-C=N feature at 4.62 pm (Chiar et al. 2002). The solid CO to HzO-ice ratio is similar to that observed in local molecular clouds. In the solar neighborhood, the X-C-N feature is seen only toward some deeply embedded protostars. Toward the GC, it may indicate the serendipitous presence of such an object in the line of sight to IRS 19, or it might conceivably arise from the processing of ices in the circumnuclear ring of the GC itself. rn'
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Fig. 3 Comparison of the 3 pm ice feature observed in the local Taurus molecular cloud (toward the background star, Elias 16; Smith et al. 1993; dashed line) and in the molecular clouds along the line of sight toward the Galactic center (KO-SWS spectrum from Chiar et al. 2000; solid line). The absorption feature centered at 3.4 pm in the GC spectrum is indicative of aliphatic hydrocarbons in the diffuse ISM (see text).
0.5 NU
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Fig. 4 The two classes of hydrocarbons thought to be present in the diffuse interstellar medium. Aliphatic (chain-like) hydrocarbons exhibit C-H stretching vibration modes near 3.4 pm, whereas the C-H stretching vibration of aromatic (ringlike) hydrocarbons is centered near 3.3 pm. Adapted from Pendleton & Allamandola 2002
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3 Diffuse Interstellar Medium Absorption Features The structured 3.4 pm absorption feature (e.g., Fig. 5 , left panels) has been the focus of many laboratory investigations into the exact nature of the hydrocarbon material (see Pendleton & Allamandola 2002 for a review). While many laboratory analog materials have provided insight into the carrier of the interstellar band based on absorption signatures at 3.4 pm, longer wavelength spectroscopy obtained from space using ISO’s SWS revealed vital information regarding the corresponding deformation modes at 6.85 and 7.25 pm toward the GC (Chiar et al. 2000). The relative strengths of these three features (Chiar et al. 2000), along with a detailed analysis of laboratory data produced via competing processes, have revealed that hydrogenated amorphous carbon produced through plasma processing, closely matches the interstellar data (Pendleton & Allamandola 2002). Our UKIRT-CGS4 spectroscopy of the GC central cluster sources, described i n section 2, shows that the depth of the aliphatic hydrocarbon absorption at 3.4 p m varies by a factor of I .7, indicating significant changes in the foreground extinction across the small field. Our spectroscopy of multiple sightlines allowed for deconvolution of the diffuse ISM absorption from the dense cloud absorption component (Chiar etal. 2002). Many previous studies of the 3.4 pm feature relied on fitting a local continuum over the 3.3 to 3.7 pm region. Our method helped us uncover broad absorption on the blue shoulder of the 3.4 pm absorption feature in the approximate spectral region where polycyclic aromatic hydrocarbons (PAHs; Fig. 4, left insert) are expected to absorb. The depth of the absorption does not correlate with the depth of the ice band, thus it is a characteristic of the diffuse ISM dust along the line of sight (Chiar etal. 2002). However, the width of the “feature” is too great to be simply reconciled with PAHs in solid grain material. The nature of this broad absorption is still an open question. Our combined UKIRT-CGS4 and ISO-SWS spectroscopy in the 3 and 6 p i regions revealed significant differences between the spectra of the GC central sources and the Quintuplet sources (Fig. 5 ) . For instance, in addition to the broad shoulder described above, a distinct narrow 3.28 Irm absorption feature is present in the Quintuplet cluster spectra (Chiar et al., in preparation), but is (probably) absent in the spectra of the GC central cluster sources. The central wavelength and width of the absorption feature are indicative of the C-H stretching vibration in PAHs. Whether the carrier of the absorption is intrinsic to the Quintuplet cluster sources or a widespread diffuse ISM dust component is unclear. The 6 pm spectra of the Quintuplet-proper sources exhibit an absorption feature centered at 6.2 pm, markedly different than the symmetric absorption feature present in the GC IRS 7 spectrum. We discuss a possible explanation for these profile differences below.
4 The Enigmatic Quintuplet Sources The five bright Quintuplet sources were discovered in near-IR surveys by Okuda et al. (1990) and Nagata ct al. (1990). We will refer to these original sources as the Quintuplet-proper sources. The surrounding cluster was revealed by later surveys (e.g., Moneti et al. 1992), and most recently by the Hubble Space Telescope which revealed hundreds of sources (Figer et al. 1999). Some of these stars have been classified as Wolf-Rayet stars, Luminous Blue Variables and OB supergiants (Figer, McLean, & Morris 1999). However, the nature of the Quintuplet-proper sources remains uncertain due to non-detection of pholospheric features which would enable their spectral classification. Each of the proposed identifications - massive dust-enshrouded young stars, OH-IR stars, dusty late-type carbon-class (WC) Wolf-Rayet stars - has its problems (Nagata etal. 1990; Figer, McLean, & Moms 1999; Moneti etal. 2001). We favor the dusty late-type WC star hypothesis. Figer, McLean, & Moms (1999) were the first to suggest that the Quintuplet sources may be dusty late-type WC stars. Moneti etal. (2001) analyze the spectral energy distributions (SEDs) of the Quintuplet sources and find that they are best reproduced by disk models similar to those used by Williams et al. ( I 987) to model SEDs of dusty late-type WC stars. However, the lack of the nearIR emission lines in the Quintuplet spectra normally used to classify the WC stars is problematic (Figer, McLean, & Morris 1999).
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Fig. 5 The 3 and 6 p m spectral regions for the Quintuplet-proper members, the GC central cluster sources, and a dusty late-type carbon-class Wolf-Rayet star (WR 118). Spectra are from Chiar et al. 2003 (Quintuplet-proper,3 pm; UKIRT-CGS4), Chiar et al. 2000 (Quintuplet (GCS3) and GC Central, 6 pm; ISO-SWS), Pendleton et al. 1994 (WR 118, 3 pm; Infrared Telescope Facility), Chiar & Tielens 2001 (WR 1 18, 6 p m ; [SO-SWS), Chiar et al. 2002 (GC Central, 3 pm; UKIRT-CGS4). The 3 pm Quintuplet spectrum is an average of the individual spectra centered on each of Quintuplet-proper members. The GC Central 3 pm spectrum is an average of 5 spectra, each has been scaled to T = 1.0 near 3.4 pm. The 6 pm spectrum of GCS3 was centered on GCS3-I; the ISO-SWS beam included GCS 3-11 and GCS3-I11(partially).
Our 6 pm spectroscopy reveals a similarity between the Quintuplet proper sources and dusty WC stars (Fig. 5). A distinct 6.2 pm absorption feature is seen toward several dusty WC stars (Schutte et al. 1998; Chiar & Tielens 2001) and the Quintuplet proper sources. The absorption feature has been attributed to the C-C stretch in amorphous carbon in the hydrogen-deficient circumstellar material associated with the WC stars, rather than PAHs in the interstellar dust along the line of sight (Chiar & Tielens 2001). Fig. 5 displays the 6 pm spectra from ISO-SWS of the lines of sight toward the Quintuplet source GCS 3-1, GC IRS 7, and the WC star WR 118. We note that the 3.4 pm hydrocarbon feature observed toward WR 1 18 is carried by the 12 magnitudes of interstellar visual extinction along the line of sight, and is not circumstellar in nature. Due to the extreme hydrogen deficiency in the WC star circumstellar environment, it is not possible to form hydrocarbon material. A broad symmetric 6.0 pm absorption feature is seen in the GC IRS 7 spectrum; the absorption is mostly accounted for by ices in the dense cloud material along the line of sight. The similarity between the WC star spectra and the GCS 3 spectrum and the dissimilarity of the GC IRS 7 spectrum is striking. These spectra are the first hint of some spectroscopic similarity between the Quintuplet sources and any of the proposed classifications, and lends support to the suggestion that they are dusty WC stars.
5
Conclusions and Future Work
Our recent spectroscopy of lines of sight toward the GC central cluster and the Quintuplet cluster has given us great insight into the chemical composition, characteristic absorption profiles, and distribution of both the diffuse ISM and dense clouds components along the line of sight. In addition, our group has carried out a program of narrow-band imaging in order to fully map the variation of the ice, hydrocarbon, and silicate dust components toward the GC central cluster, including the circumnuclear ring (Adamson et al. 2003). Future spectroscopic observations from airborne (SOFIA) and ground-based observatories will answer such questions as 1. Do icy grain mantles in the circumnuclear ring contribute to the deep ice features and
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X-CEN feature observed in those lines of sight'? 2. Is the 6 pm absorption profile of the Quintuplet-proper members unique to those sources and dusty late-type WC stars? 3. Is the distinct 3.28 pm feature observed in the Quintuplet-proper spectra due to PAH absorption in the diffuse ISM? 4. Will this feature b e present in high signal-to-noise spectra of other diffuse ISM lines of sight such as heavily extincted WC stars? 5. Is this feature really absent in lines of sight toward the GC central cluster? T h e answers to questions 2 through 5 will not only tell us more about aromatic hydrocarbons in the diffuse ISM, but will also bring us closer to uncovering the nature of the enigmatic Quintuplet sources. Acknowledgements J.E. Chiar gratefully acknowledges support from NASA's Long Term Space Astrophysics Program under grant 399-20-61-02.
References Adamson, A. J., et al., these proceedings. Adamson, A. J., Whittet, D. C. B., Chrysostomou, A., Hough, J. H., Aitken, D. K., Wright, G. S., & Roche, P. F. 1999, ApJ, 512,224 Allamandola, L. J., Tielens, A.G. G. M., & Barker, J. R. 1989, ApJS, 71,733 Boogert, A. C. A., Helmich, F. P., van Dishoeck, E. F., Schutte, W. A,, Tielens, A. G. G. M., & Whittet, D. C. B. 1998, A&A, 336,352 Brooke, T. Y., Sellgren, K., & Smith, R. G. 1996, ApJ, 459,209 Butchart, I., McFadzean, A. D., Whittet, D. C. B., Geballe, T. R., & Greenberg, J. M. 1986, A&A, 154, L5 Chiar, J. E., Adamson, A. J., & Whittet, D. C. B. 1996, ApJ, 472,665 Chiar, I. E., Pendleton, Y. J., and Geballe, T. G., and Tielens, A. G. G. M. 1998, ApJ, 507,281 Chiar, J. E., & Tielens, A. G. G. M. 2001, ApJ, 550, L207 Chiar, J. E., Tielens, A. G. G. M., Whittet, D. C. B.. Schutte, W. A,, Boogert, A. C. A., Lutz, D., van Dishoeck, E. F., & Bernstein, M. P. 2000, ApJ, 537, 749 Chiar, J. E., Adamson, A. J., Pendleton, Y. J., Whittet, D. C. B., Caldwell, D. A,, & Gibb, E. L. 2002, ApJ, 570, 198 Dartois, E., & d'Hendecourt, L. 2001, A&A, 365, 144 de Graauw, T., Whittet, D. C. B., Gerakines, P. A., et al. 1996, A&A, 3 IS, L345 Figer, D. F., McLean, I. S., & Moms, M. 1999, ApJ, 514,202 Figer, D. F., Kim, S. S., Moms, M., Serabyn, E., and Rich, R. M., & McLean, I. S. 1999, ApJ, 525,750 Geballe, T. R., Baas, F., & Wade, R. 1989, A&A208, 255 Gerakines, P. A,, Whittet, D. C. B., Ehrenfreund, P., er al. 1999, ApJ, 522, 357 Gusten, R.,Genzel, R., Wright, M. C. H., Jaffe, D. T., Stutzki, J.. & Harris, A. 1987, in AIP Conf. Proc. 155: The Galactic Center, ed. D. C. Backer (New York, AIP). 103 Ishii, M., Nag_ata,T., Chrysostomou, A., & Hough, J. H. 2002, AJ, 124, 2790 Lebofsky, M.J. 1979, AJ, 84, 324 Li, A., & Greenberg, J. M. 2002, ApJ, 577, 789 Lutz, D., Feuchtgruber, H., Genzel, R., et al. 1996, A&A, 3 15, L269 McFadzean, A. D., Whittet, D. C. B., Bode, M. F., Adamson, A. J., & Longmore, A. J. 1989, MNRAS, 241,873 Moneti, A., Glass, I., & Moorwood, A. 1992, MNRAS, 258,705 Nagata, T., Woodward, C. E., Shure, M., Pipher, J. L., & Okuda, H. 1990, ApJ, 351, 83 Okuda, H. et al. 1990, ApJ, 351, 89 Pendleton, Y. J., & Allamandola, L. J. 2002, ApJS, 138, 75 Pendleton, Y. J.. Sandford, S. A., Allamandola. L. J . , and Tielens, A. G . G. M., & Sellgren, K. 1994, ApJ, 437,683 Sanders, D. B., and Scoville, N. Z . , & Solomon, P. M. 1985, ApJ, 289,373 Sandford, S. A., Allamandola, L. J., Tielens, A.G.G.M., Sellgren, K., Tapia, M., &Pendleton, Y.J. 1991, ApJ, 371, 607 Schutte, W. A., van der Hucht, K. A., Whittei, D. C. B., et al. 1998, A&A, 337,261 Smith, R. C., and Sellgren, K., & Brooke, T. Y. 1993, MNRAS, 263,749 Tielens, A. G. G. M., Wooden, D. H., Allamandola, L. J., Bregman, J., & Wittebom, F. C. 1996, ApJ, 461,210 Whittet, D. C. B., Boogert, A. C. A., Gerakines, P. A., et al. 1997, ApJ, 490, 729 Williams, P. M., van der Hucht, K. A., & The, P. S. 1987, A&AI82, 91 Willner, S. P., Russell, R. W., Puetter, R. C., Soifer, B. T., & Harvey, P. M. 1979, ApJ, 229, L65 Wright, G., Geballe, T., Bridger, A., & Pendleton, Y.J. 1996, in New Extragalactic Perspectives in the New South Africa, eds. D.L. Block & J.M. Greenbeg (Dordrecht:Kluwer), 143
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Astron. Nachr./AN 324, No. SI, 117- 123 (2003) / DO1 10.1002/asna.20038S064
X-rays from the HI1 Regions and Molecular Clouds near the Galactic Center Katsuji Koyama* I , Hiroshi Murakami** 2 , and Shinichiro Takagi***
'
' Department of Physics, Kyoto University, Kyoto 606-8502, Japan Institute of Space and Astronautical Science(ISAS), Kanagawa 229-85 10, Japan
Key words X-rays, Molecular clouds, HI1 regions, Young stars
Abstract. We report measurements by C,'hnn,dra of a variety of X-ray sources in the molecular clouds and HI1 regions of the Sgr B2, Arches, Quintuplet and the Galactic center clusters. Moderately bright X-ray sources are present in the Sgr B2, Quintuplet and the Galactic center clusters at the positions of ultra compact HI1 regions and bright infrared sources. Their X-ray spectra are fitted with models of a thin thermal plasma with 2-10 keV temperatures and luminosities of 10"2p""erg s-'. The X-ray properties are typical of those of high-mass young stellar objecta or clusters of such objects. The Arches Cluster has three bright X-ray sources, at the positions of bright IR and radio stars, with X-ray luminosities of a few x l0""erg spl each, which may indicate an unusual X-ray emission mechanism from high mass YSOs. A unique X-ray feature of molecular clouds and HI1 regions near the Galactic center is the presence of diffuse emission with a strong 6.4 keV line; in Sgr B2 this is attributable to the fluorescence of gas irradiated hy external sources in the Galactic center, while the diffuse emission from Arches is puzzling.
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1 Introduction A unique X-ray feature in the Galactic center (GC) region was discovered by the Ginga satellite. The X-ray spectrum near the GC exhibits a bright iron line (Koyama et al. 1989). Figure 1 shows scan profiles taken with the classical proportional counter of beam size about I degree. The top panel is the wide hand X-ray flux (2-18 keV), and there is no large peak at the GC. However, a narrow band X-ray profile that includes the 6.7 keV iron line shows a bright peak at the GC, by far the brightest source (middle panel). The emission near 6.7 keV is extended and positive longitudes (north-east of the GC) emit lower energy (about 6.6 keV) than do negative longitudes, which emit at 6.7 keV (see lower panel of figure 1). In order to determine the origins of the iron line emission and the energy variations, we have used the ASCA satellite to obtain imaging spectroscopy of the GC region. Figure 2 (left) shows the ASCA image of the GC region in the 2-10 keV band. As expected, we found bright diffuse X-rays near the GC extending over a 1 square degree area. Figure 3 is the X-ray spectrum of this diffuse component. It shows many emission lines from highly ionized Si, S, Ar, Ca and Fe. The spectrum implies that the origin of the diffuse X-ray is a thin hot plasma of 107-108K. The continuum X-rays accurately trace this high temperature plasma. A surprising discovery is that there is a 6.4 keV line just below the 6.7 and 6.9 keV lines. The latter two lines are due to He- and H-like irons, and hence support the presence of very high temperature plasma (near IO'K), while the former line is due to neutral or low ionization states of iron. To investigate t h e origin of the 6.4 keV line, we obtained a narrow band image, which is shown in figure 2 (right panel). The
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* e-mail: koyama 0crscphys.kyoto-u.ac.jp * * e-mail:
[email protected] * * * e-mail: takagi9&cr.scphys.kyoto-u.ac.jp
@ 2003 WILEY-VCH Verlag GnihH & Co KGaA, Wcinhrim
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Fig. 1 The scan profile of the X-ray emission in the 2-18-keV energy band ( u: upper panel), the iron-line emission (b: middle), and the line energy along the Galactic plane (c: lower panel), (from Koyarna et al. 1989).
Fig. 2 The ASCA images of the GC region. Courtesy: Yoshitomo Maeda (see http://www-cr.scphys.kyoto-u.ac.jplIAUlgallely/gallery.html)
figure shows that the 6.4 keV line emission is clumpy, and is more concentrated to the northeast of the G C than to the opposite side (south west). This asymmetry is the origin of the Ginga asymmetry of the iron line energy.
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Fig. 3 The X-ray spectrum of the GC region (from Koyama et al. 1996)
By comparing the X-ray and radio data, we found that the 6.4 keV line emissions are well correlated with the giant molecular clouds (MC), which are high mass star forming regions. The region of brightest 6.4 keV line emission is the Sgr B2 cloud, about 0.65 degree (or lOOpc in projection) from the GC. The next brightest region is just inside of the Radio Arc, which corresponds to a newly found molecular cloud (MC) MO. I3+0.13. There also is faint emission near the Arches cluster, which contains young high mass stars and many compact HI1 regions. In the flowing sections, we report more details of individual clouds, in particular the X-ray emission from high mass young stellar objects (YSO) and the diffuse 6.4 keV line emission. Since MCs may harbor high mass YSOs, we briefly review the X-ray emission from high mass YSOs observed with Chnndl-a. Although the samples are still limited, the X-ray luminosities of YSOs appear to increase with increasing stellar mass, and saturate at about l0"'org s-' for massive stars or stars associated with compact HI1 regions (e.g. Kohno et al. 2002). The X-ray spectra of high mass YSOs are generally harder (a few keV) than stars on the main sequence (less than 1 keV).
2
SgrB2
Although ASCA found diffuse X-rays with strong 6.4 keV line emission from Sgr B2, X-rays from many unresolved point sources (possibly YSOs) may contaminate the diffuse emission. We thus made high spatial resolution observation of Sgr B2 with Chandra (ObsID = 944), and found 17 new point sources in the cloud. Figure 4 is the C h m d r a image near the center of Sgr B2 overlaid on the radio contour map. The two brightest X-ray sources (Nos. 10 and 13) lie near Sgr 8 2 Main (M), which contains the largest complex of ultra compact HI1 (UCHII) regions Figure 5 is a close-up view of Sgr B2 (M). The X-ray emission originates from several point sources, some of which coincide with ultra compact radio sources, while others have no radio counterpart. The
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composite X-ray spectra near from the UCHII complexes 10 and 13 are fitted with models of thin thermal plasmas with temperatures of 5-10 keV and large absorption column densities of 4 x H cm-2. The high temperatures are consistent with the high-mass YSOs in the Orion Nebula and the Mon R2 cloud (e.g., Shultz et al. 2001; Kohno et al. 2002). The X-ray luminosities from these UCHII complexes are about 10"crg spl. Since the UCHII complexes are typically comprised of about 10 sources, the X-ray luminosity of each point source is about 103'erg s-', which is roughly consistent with high mass YSOs. These luminosities, together with the high H indicate that some of these hard X-ray sources are absorption column densities of 4 x likely due to high-mass YSOs in the core of the Sgr B2 (M). Others are UCHII regions, and hence are probably zero-age main sequence (ZAMS) stars. N
The other UCHII complexes, Sgr B2 North (N) and South (S), also show hard X-rays with high abrorpH cmP2 and luminosities of a few to 10 times 10"crg s-'. tion column densities of about (5-6) x Thus these are also clusters of high mass YSOs located in the cores of the UCHII complexes. We note that the absorption column density of 5 x loz3H cm-2 for the Sgr B2 X-rays is the largest among the known stellar X-ray sources. Our observations demonstrate that hard X-rays are a very effective way to discover deeply embedded, YSOs even those that are located as far away as the GC and suffer from optical extinctions as large as of A v -200-300 mag. N
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Fig. 4 The 2-10 keV band X-ray image of Sgr B2, overlaid on the contours of HI1 regions (Gaume et al.
Fig. 5 The 2-10 keV band X-ray image near Sgr 8 2 (M), overlaid on the contours of HI1 regions (De Pree et al. 1998, Gaume & Claussen 1990). The right source is No. 10, and the left source is No. 13.
1995).
In addition to the point sources, we also detected diffuse X-rays in the 6.4 keV line. The integrated flux of point sources (YSOs) is only 10%of the diffuse flux. Therefore at 6.4 keV the diffuse X-rays dominate the total emission. The morphology of the 6.4 keV line emission is a crescent with the curvature pointing toward the GC. This morphology, together with the spectrum, can be reproduced by illumination by strong X-rays from the GC direction. We call this class of X-ray source an "X-ray Reflection Nebula (XRN)."
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3 Arches Cluster In the Chnrrdru mosaic image of the GC region (obsID= 242+945+1561), there is also a complex of Xray emission from two HI1 regions, the Arches and Quintuplet clusters. Figure 6 is a close-up view of the Arches cluster. Three of the bright X-ray sources coincide with the positions of bright IR andlor radio sources. We fitted the X-ray spectra for the 3 sources with a thin thermal plasma model with solar metal abundances. The best-fit parameters are given in Table I . The temperature of 1-3 keV and the large X-ray H ciiir2 are consistent with these sources being high mass YSOs absorption column density of about in the cluster. The X-ray luminosities (in the 2-10 keV band) are, however, a few x 10””t.i-gsrl, which are the highest among known YSOs. They are similar to those of the 30 Dor YSOs, which possibly are high-mass YSO binaries producing strong X-rays in the collisions of energetic stellar winds.
2
1
R3
Fig. 6 The 2-10 keV band X-ray iinage of the Arches cluster. The bin size is 0.”5 x 0.”5. The diamond ( 0 )and cross (+) show the position of the bright infrared sources (Blurn et al. 2001) and the radio sources (Lang et al. 2001), respectively.
Table 1 The best-fit parameters for the YSOs in the Arches Clustei
ID:IR(Radio)
26
2 1 (AR4)
23(AR1) _____
_____
N H( 10”Hcrr~-2)
6.9(4.5-11.2)
9.3(5.6-15.5)
9.9(6.5-14.8)
kT(keV)*
2.4(1.3-5.6)
1.3(0.6-2.5)
1.6(1.0-2.7)
L,(lO”erg s-l)
2.8(2.3-3.4)
3.9(3.14.7)
5.0(4.2-6.0)
* in the 2-l0keV
band
Surrounding the Arches Cluster is a diffuse structure (a circle of about 1’ radius), which shows strong 6.4 keV line emission with deep absorption at low energy (Yusef-Zadeh et al. 2002). The diffuse spectrum is fitted with a power-law and a narrow Gaussian line at 6.4 keV. This spectral feature is similar to Sgr B2, an XRN. However the X-ray morphology is different from Sgr B2, and is roughly a circle around the three brightest YSOs. Therefore, unlike Sgr B2, the irradiating source may not be external. YSOs within the
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Arches cluster may be the sources. For simplicity, we assume an uniform density spherical cloud, of size similar to the diffuse emission of 50" radius (2.1 pc at 8 kpc). From the observed NH the mean density Using the total flux from YSOs (the brightest three sources), we of the cloud is estimated to be 10'c&'. calculate the expected diffuse 6.4 keV line luminosity to be 1.3 x 10"erg s-', which is only 10% of the 103'erg s-l. Yusef-Zadeh et al. (2002) proposed an exotic scenario in which the observed value of diffuse X-rays are shock-heated gas created by the collision of individual 1000 km/s stellar winds in the dense cluster environment. However this hot gas should emit the 6.7 keV line, not the 6.4 keV line. Thus the origin of the 6.4 keV line is a puzzle and a deeper exposure Cliaridra observation is needed to help resolve it.
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4 The Quintuplet Cluster The Quintuplet cluster is one of the most massive clusters in the Galaxy and contains bright IR sources and radio emission from the HI1 region (e.g., Figer et al. 1999 and references therein). Most of the cluster sources are thought to be 09 and WR stars (Figer et al. 1999); hence X-rays from high-mass YSOs are expected from this region. In the 50-ksec pointing observation (obsID= 945), the Quintuplet is located in a gap between CCD chips. Therefore, we have examined the GC survey data (obsID=2276, 12-ksec exposure). Although there are excess X-rays from this region, the limited exposure time makes it difficult to resolve individual sources. Assuming the same temperature and absorption as those of the Arches cluster, we estimate their luminosities to be 4 x 10"erg s-', which may be typical of high-mass YSO clusters. To make further progress, a deeper exposure is needed.
5 The Galactic Center Clusters In the GC itself X-rays are found from some of the IR bright stars (or star clusters), IRS 3, 13, 31 and 16SW (Baganoff et al. 2001). The GC region is known to be a massive star-forming region (e.g., Morris 1993; Paumard et al. 2001), and these IR sources are likely high mass stars. For example, IRS 13 is known to consist of many stars including the WR star IRS 13E (Paumard et al. 2001). IRS 16SW is probably an eclipsing He-star binary, hence the X-rays are probably due to the colliding stcllar winds (Ott et al. 1999). Therefore the observed X-ray sources are likely individual YSOs or clusters. Their average absorption column densities of 1023Hc1n-l and luminosities of 103't5rg s-' (Baganoff et al. 2001) are typical of high-mass YSO clusters in the GC region.
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6 Summary We report C h a ~ ~ d robservations a. of the Sgr B2, Arches, Quintuplet and the Galactic center clusters. Sgr 9 2 exhibits X-ray emission coincident with three UCHII complexes (Main North and South). The X-ray spectra are fitted with thin thermal plasma models of 5-10 keV temperature, which are consistent with the sources being high-mass YSOs. Diffuse emission with strong 6.4 kcV lines is round, which is most likely due to external irradiation from the direction of the GC. The Arches Cluster contains three bright X-ray sources, at positions of bright IR and radio sources, with X-ray luminosities of a few x l0"erg 8 - l (2-10 keV band) each: these are the highest among any known YSOs. There is also diffuse emission, which includcs strong 6.4 keV line emission. Unlike Sgr 92, the 1' radius, suggesting that the irradiating sources are internal. However the morphology is a circle of YSOs can produce only 10% of the observed line flux, hence the true origin is still puzzling. There also are excess X-rays from the Quintuplet cluster. Assuming the same temperature and absorption as those of the Arches cluster, we estimate the luminosity to he 4 x 103'erg s-l, which may be typical of high-mass YSO clusters. N
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References Baganoff, F. K. , Maedd, Y. , Moms, M. , Bautz, M. W., Brandt, W. N.,Cui, W., Doty, J. P. , Feigelson, E. D . , Garmire, G. P. , Pravdo, S. H. , Ricker, G. R. ,& Townsley, L. K. 2001, astro-ph/0102151 Blum, R. D., Schaerer, D., Pasquali, A., Heydari-Malayeri, M., Conti, P. S., & Schmutz, W. 2001, AJ, 122, 1875 De Pree, C. G., Goss, W. M., & Gaume, R. A. 1998, ApJ, 500,847 Figer, D. F., McLean, I. S., & Morris, M. 1999, ApJ, 5 14,202 Gaume, R. A. & Claussen, M. J. 1990, ApJ, 351, 538 Gaume, R. A., Claussen, M. J., De Pree, C. G.. Goss, W. M., & Mehringer, D. M. 1995, ApJ, 449,663 Kohno, M., Koyama, K., & Hamaguchi, K. 2002, ApJ, 567,423 Koyama, K., Awaki, H., Kunieda, H., Takano, S., & Tawara, Y. 1989, Nature, 339, 603 Koyama, K., Maeda, Y., Sonobe, T., Takeshima, T., Tanaka, Y., & Yamauchi, S. 1996, PASJ, 48,249 Lang, C. C., Goss, W. M., & Rodriguez. L. E 2001, ApJL, 551, L143 Morris, M. 1993, ApJ, 408,496 Murakami, H., Koyama, K., I%Maeda, Y. 2001, ApJ, 558, 687 Murakami, H., Koyama, K., Tsujimoto. M., Maeda, Y., & Sakano, M. 2001, ApJ, 550, 297 Ott, T., Eckart, A., & Genzel, R. 1999, ApJ, 523, 248 Paumard, T., Maillard, J. P., Moms, M., & Rigaut, F. 2001, A&A, 366, 466 Senda, A., Murakami, H., & Koyama, K. 2002, ApJ, 565, 1017 Schulz, N.S., Canizares, C., Huenemoerder, D., Kastner, J. H., Taylor, S. C., & Bergtrom, E. J. 2001, ApJ, 549,441 Takagi, S., Murakami, H., & Koyama, K. 2002, ApJ, 573, 275 Yusef-Zadeh, F., Law, C., Wardle, M., Wang, Q. D., Fruscione, A,, Lang, C. C., & Cotera, A. 2002, ApJ, 570, 665
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Astron. NachrJAN324, No. SI, 125- 131 (2003) I DO1 10.1002/asna.200385075
Reflected X-ray Emissions on Giant Molecular Clouds -Evidence of the Past Activities of Sgr A* Hiroshi Murakami*’, Atsushi Sends**?, Yoshitomo Maeda***’,and Katsuji Koyamat’
’ High Energy Astrophysics Division, Institute of Space and Astronautical Science 3-1 -1 Yoshinodai, Sagamihara, Kanagawa 229-8510, Japan ’ Department of Physics, Graduate School of Science, Kyoto University Kitashirakawa, Sakyo-ku, Kyoto, 606-8502, Japan
Key words molecular clouds, X-ray, refection, Sgr A*, Sgr B2, Sgr C, Radio Arc, MO.ll-0.08 Abstract. We have found strong 6.4-keV line emissions from the giant molecular clouds in the Galactic center region: Sgr B2, Sgr C , and M0.I 1-0.08 (at the Radio Arc region). The high angular resolution of Chcrndra reveals that the 6.4-keV line emissions are indeed coincident with the clouds, and shifted towards the Galactic center. The X-ray spectra have very strong 6.4-keV lines with equivalent widths 2 1 keV and are attenuated by larger column densities of interstellar gas. These characteristics imply that the massive molecular clouds are irradiated by an external X-ray source in the direction of the Galactic center and emit fluorescent and scattered X-rays. These clouds are new category of X-ray source: “X-ray Reflection Nebula”. The reflected X-ray flux reveals the recent luminosity history of the primary irradiating source, which may be the massive black hole Sgr A*, according to the light travel time to each cloud. Making use of the radio determinations of the cloud masses, we find that Sgr A* was as luminous as 10” erg s-’ a few hundred years ago, and has gradually decreased to present value.
1 Introduction With its wide energy X-ray band (0.5-10 keV) imaging capability, ASCA found 6.4-keV line emissions in the Galactic center (GC) region (Figure 1; Koyama et al. 1996). There are two distinct peaks, the Sgr B2 region and the Radio Arc region, which roughly agree with the locations of giant molecular clouds. For Sgr B2, by comparison with a radio observation of the molecular cloud, we have found that the distribution of 6.4-keV line is indeed correlated with the cloud, and is slightly shifted from the cloud core (Koyama et al. 1996;Murakami et al. 2000). Because the 6.4-keV line is characteristic of radiation from neutral iron, it i s natural to think that X-rays are emitted by molecular clouds. However, the clouds are very cold, and cannot emit high energy X-rays in themselves. Koyama et al. (1996) suggested that the clouds are irradiated by an external X-rays, and emit fluorescent X-rays in the 6.4-keV line. Thus the Sgr B2 cloud may b e a new category of X-ray source, an X-ray Reflection Nebula (XRN). To explain the luminosity of 6.4-keV line from the Sgr B2 cloud, a strong X-ray source with the luminosity of ,-- lo3’ erg sP1 is required (Murakami et al. 2000), but there is no bright X-ray source in the G C region. The most plausible explanation is that our Galactic nucleus Sgr A*, which is a massive black hole * e-mail: hiroQastro.isas.ac.jp,Phone: +81 427598134, * * e-mail: sendaQcr.scphys.kyoto-u.ac.jp *** e-mail: ymaedaQastro.isas.ac.jp t e-mail: koyamaQcr.scphys.kyoto-u.ac.jp
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(MBH) of 1O6A4a,was much brighter some hundreds of years ago, the light travel time between Sgr A* and Sgr B2, and is dim at present. N
The limited spatial resolution of ASCA, however, could not exclude possible contamination of many point sources such as young stellar objects in Sgr B2, which might deform the spectrum and morphology of the diffuse emission. Hence we have analyzed Chandra observations of giant molecular clouds in the GC region: Sgr B2, Sgr C, and M0.11-0.08 (at the Radio Arc region). With the higher angular resolution of Chandru, we have been able to constrain the contribution of point sources to a negligible level. Throughout this paper, the distances to the clouds are assumed to be 8.5 kpc, the same as to Sgr A*.
Fig. 1 The distribution of the 6.4-keV line intensity in the GC region observed with ASCA (Koyama et al. 1996).
2 Observations
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Chandra observed Sgr B2 and MO. 11-0.08 with long exposure times of 100 ksec (ObsID=944) and 50 ksec (ObsID=945), respectively. The Sgr C region was observed as a part of GC survey, with a net exposure time of 20 ksec (ObsID=2267+2270+2272). We screened the observed data using standard criteria, and made point source subtracted spectra and 6.4-keV line images. N
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3 Results 3.1
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Sgr B2 is one of the most massive giant molecular clouds in the Galactic center region, and the most luminous in the 6.4-keV line band. With Cliandm's high angular resolution, we detected 20 point sources in the cloud. Some of the sources are coincident with star forming regions, and are considered to be young stellar objects. The integrated luminosity of all the resolved point sources is 3 x 10"" erg s - l , which is only 3% of the luminosity of the diffuse X-rays (see below). The details for the point sources are given in a separate paper (Takagi et al. 2002). Figure 2a shows the 6.4-keV line image of Sgr B2. In this energy band most of the X-rays are from diffuse emission. The diffuse X-rays come mainly from the south-west half of the cloud and the emission region has a concave shape pointing towards the Galactic center. The point source subtracted spectrum (Figure 3 ) exhibits strong Kcu and K,G emission lines of neutral iron. The center energy of K u line (Ec;) is 6.38 keV, and the feature has an equivalent width (EW) of 1.7 keV. The continuum component can be reproduced by a power-law model with a fixed photon index of 2.0 and a large absorption column density of NH 8.8 x cmp2. The luminosity is about erg spl (4-10 keV).
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Fig. 2 (a) 6.4-keV line image of Sgr B2 observed with Chandru. The contours show the density distribution of the molecular cloud ("CO, Sato et al. 2000). The point source at the center of the cloud is coincident with the star forming region Sgr B2 Main (Takagi et al. 2002). (hj A simulated image with the X-ray retlcction nebula model. The irradiating source is assumed to be in the direction of CC.
To confirm the XRN scenario, we made a numerical simulation of the reflected X-rays for the casc of the molecular cloud being irradiated by an external X-ray source in the direction of the GC. The simulated image is shown in Figure 2b. The shifted distribution and the crescent shape are well reproduced by the simulation. The simulated spectrum with the XRN model exhibits very strong 6.4-keV line and a large absorption, and fits the observed spectrum well, with a reduced x2 of 1.14 (solid line in Figure 3).
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These results support the idea that the Sgr B2 emission is fluorescent and scattered X-rays due to irradiation by an external source in the direction of the GC. The details for Sgr B2 cloud and XRN simulations are already published in Murakami et al. (2001).
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Fig. 3 The spectrum of Sgr B2. The solid line shows the simulated spectrum with an XRN model (Murakami et al. 2001), and the dashed line shows the contribution of point sources.
3.2 SgrC This cloud is located at the opposite side of the GC from Sgr B2. The 6.4-keV line image (Figure 4a) shows diffuseX-ray emission on the GC side (at the upper left) of the cloud. The distribution has a crescent shape, which is oriented similarly with respect to the GC as that of Sgr B2. The spectrum (Figure 4b) is well fitted with the same model as Sgr B2, and exhibits strong iron line (Ec 6.38 keV; EW 3.2 keV) and a large absorption ( N H 1.4 x loz3 cmp2). These features are also same as Sgr B2. Thus this cloud also is a candidate XRN. More detailed analysis, such as comparison with a numerical simulation, is meaningless because of the poor photon statistics. An additional Chundru observation with longer exposure time is required.
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3.3 M0.11-0.08 MO. 11-0.08 is a molecular cloud near the Radio Arc, and the nearest giant molecular cloud to the GC. Figure 5a, b show the 6.4-keV line image and the spectrum, respectively, of MO.ll-0.08. Diffuse X-ray emission fills the entire region of the cloud. In addition, filamentary structure is found at the edge of the cloud facing the GC. The spectrum is roughly reproduced by a thin thermal plasma model, but there are residuals at 6.4 keV and at higher energy. Therefore, we fitted the spectrum with a thermal component and the same model as Sgr B2 and Sgr C (a power-law and two Gaussian lines of K N and KB). The best-fit parameters of the K n line are, Ec 6.36 keV and EW 1.1 keV. The absorption column density for the power-law component is also large ( N H 2.4 x loz3cm-2). The thermal component is reproduced by emission from a thin thermal at a temperature of 3 keV and having a luminosity of 3 x erg s-'. Thus we find that the X-ray spectrum of MO. 1 1-0.08 exhibits strong 6.4-keV line. Although there is also a thermal component, fluorescent X-rays must be emitted from the cloud. By using the same model as Sgr B2 and Sgr C, we are able to reproduce the spectrum. MO.ll-0.08 also must be an XRN. N
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Fig. 4 (a) The 6.4-keV line image of Sgr C superposed on the radio intensity of the CS line (radial velosity of -120 km s-' to -110 km s-'; Tsuboi, Handa. Ukita 1999). Diffuse X-ray emission is seen on the GC side of the cloud. (b) The X-ray spectrum of Sgr C. The solid line shows the best-fit model spectrum using a power-law and two Gaussian lines.
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Fig. 5 (a) The 6.4-keV line image of MO. I 1-0.08. The contours are the CS line intensity over the radial velociry range of20 kin s-' to 30 hi sC1 (Tsuboi. Handa. & Ukita 1999). (b) Thc X-ray spectrum of MO.ll-0.08. The dotted line shows the thermal component, and the dashed lines show the reflected component (a power-law with two narrow lines of neutral iron KCYand K L ~ ) .
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T h e thermal component might be related to the expansion of the cloud. T h e total energy of the plasma erg. A part of the expansion energy of 10" erg (Tsuboi, Ukita, & Handa 1997) could be transformed into thermal energy and result in the emission of X-rays.
is about 3 x
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4 Discussion In the previous section, we found that three molecular clouds are probably XRNe. If so, the luminosity of the 6.4-keV line from each cloud indicates the past luminosity of the irradiating source at a time corresponding to the light travel time from the source to the cloud. As already mentioned, the irradiating source is required to be very luminous, and considered to be the Galactic nucleus Sgr A*. We thus convert the distance from Sgr A* to elapsed time, and make a light curve for the X-ray luminosity of Sgr A* during the past 500 years (Figure 6). The luminosity was as high as lo3' erg s-' a few hundreds years ago, and seems to have decreased gradually to the present value. However, there are only four data points, and thus we cannot discuss the detailed variability. The time span of each data point corresponds to the size of each cloud. Our method is insensitive to variability on shorter time scales. The past activity of Sgr A" could have been generated by a surge of accretion onto the MBH due, for example, to the passage of dense shell of a young supernova remnant Sgr A East as discussed by Maeda et al. (2001) based on the new Chnndru results on the GC. They estimated that the age of the SNR is lo4 years, and that the dense shell reached Sgr A* lo3 years ago. After the shell swept the surrounding matter past Sgr A* the luminosity would have become anomalously low. Their time scale agrees with the derived light curve by the XRN model. Hence our drivation of the past activity of Sgr A* suggest that Sgr A* is in a low luminosity phase at present.
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Time (year) Fig. 6 The past luminosity of Sgr A* estimated from the luminosity of 6.4-keV line of each cloud.
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Summary
We observed three molecular clouds, Sgr B2, Sgr C, and M0.11-0.08 with Chundvu. With its high angular resolution, we found diffuse neutral iron line emissions from all of these clouds. These giant molecular clouds are "X-ray Reflection Nebulae".
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From the intensity of reflected X-rays, we obtained the luminosity history of the Galactic nuclei Sgr A* during the last 500 years. Sgr A* was as luminous as erg s-l a few hundreds years ago, and has dimmed gradually since then
References Koyama, K., Maeda, Y., Sonohe, T., Takeshima, T., Tdnaka, Y., & Yamauchi, S. 1996, PASJ, 48, 249 Maeda, Y. et al. 2002, ApJ, 570,671 Murakami, H., Koyama, K., Sakano, M., Tsujimoto, M., & Maeda, Y. 2000, ApJ, 534, 283 Murakami, H., Koyama, K., & Maeda, Y. 2001, ApJ, 558, 687 Sato, F., Hasegawa, T., Whiteoak, J. B., & Miyawaki, R. 2000, ApJ, 535, 857 Takagi, S., Murakami, H., & Koyama, K. 2002, ApJ, 573,275 Tsuhoi, M., Handa, T., & Ukita, N. 1999, ApJS, 120, 1 Tsuhoi, M., Ukita, N., & Handa, T. 1997, ApJ, 48 I , 263
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Astron. Nachr./AN 324, No. S1, 133 - 137 (2001) / DO1 10.1002/asna.200385078
Observation of Toroidal Magnetic Fields on 100 pc Scales in the Galactic Center
',
G . Novak"',D. T. Chuss', J. L. Dotson 3, G. S. Griffin4, R. F. Coewenstein M. G. Newcomb 5 , D. Pernic 5 , J. B. Peterson 4, and T. Renbarger6 I
Department of Physics and Astronomy, Northwestern University
* NASA, Goddard Space Flight Center
'
NASA, Ames Research Center Physics Department, Carnegie Mellon University Yerkes Observatory,University of Chicago School of Physics and Astronomy, University of Minnesota
Key words Magnetic Fields, Galactic Center, Submillimeter Polanmetry, Interstellar Dust Abstract. We present new submillimeter polarimetric observations of the Galactic center region, made using the SPARO polarirneter that operates at the South Pole. Compared with previous submillimeter polarimetry of this region, our measurements cover much more sky area, and they imply that the molecular gas in the central few hundred pc is threaded by a large scale toroidal magnetic field. We consider this result together with radio observations that show evidence for poloidal fields in the Galactic center, and with Famday rotation observations. We compare all of these observations with a magnetodynamic model for the Galactic center.
1 Introduction Our contribution to the Galactic Center Workshop 2002 was to present new submillimeter polarimetric observations of the Galactic center, obtained using the SPARO instrument at South Pole station. SPARO (the Submillimeter Polarimeter for Antarctic Remote Observations), is a 9-pixel submillinieter array polarimeter incorporating 3He-cooled detectors (Renbarger et al. 2003). It is operated on the Viper telescope (Peterson et al. 2000). Submillimeter thermal emission from interstellar dust grains is generally polarized, due to magnetic alignment of grains. Thus, submillimeter polarimetry provides a method for mapping interstellar magnetic fields. The SPARO map extends over a much larger sky area than has been covered in previous submillimeter polarimetric maps, so it provides new information on the large-scale configuration of the magnetic field in the Galactic center. This can be compared with the results of radio synchrotron observations and Faraday rotation observations, both of which also give information about magnetic fields. We have already published a paper that presents our SPARO results, and compares them with synchrotron and Faraday observations (Novak et al. 2003). Accordingly, here we will merely summarize the Novak et al. (2003) paper.
2
SPARO results
As is typical for submillimeter continuum observations, measurements made using SPARO are not absolute, but rather are differential. Specifically, the flux that SPARO measures is the difference between the * e-mail:
[email protected],Phone: 847 491 864s
@ 2003 WILEY-VCH Vcrlilg GrnhH Rr Ca KGaA. Wetnhcioi
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flux at the main observing position and the average flux for the two sky reference positions, which are separated from the main position by +0.5” and -0.5” in cross-elevation, respectively. We observed the Galactic center for a total of five weeks during the interval April-July 2000. The results are presented in Figure 1. The contours correspond to a 450 p m photometric map made using SPARO, and they clearly show the large concentration of molecular gas that is associated with the innermost few hundred pc of the Galaxy. The highest column density occurs at the position of Sgr B2, displaced from the center of the Galaxy toward positive Galactic longitudes. The SPARO polarization results are shown using bar symbols. The orientation of each bar gives the inferred magnetic field direction, that is orthogonal to the E-vector of the measured polarization, and the length of the bar is proportional to the degree of polarization. The SPARO polarization results imply that the magnetic field permeating the Galactic center molecular gas, when projected onto the plane of the sky, is for the most part parallel to the Galactic plane. The most natural way to account for this is to suppose that the molecular gas in the Galactic center is threaded by a large scale magnetic field having a toroidal configuration. This had already been suggested, based on earlier polarimetry results at far-infraredhbmillimeter wavelengths (Morris et al. 1992, Novak et al. 2000). However, the SPARO results cover much more sky area than the previous observations, and thus they provide the strongest evidence yet obtained for the existence of this toroidal large-scale field. Figure 2 shows the SPARO magnetic “vectors” superposed on a radio continuum image (gray scale) showing Galactic center non-thermal radio filaments. These filaments trace magnetic fields running preferentially perpendicular to the Galactic plane. They appear to delineate a large scale magnetic field with a poloidal configuration. It is clear from Figure 2 that the magnetic field in the central few hundred pc is neither purely toroidal, nor purely poloidal. Rather, there appear to be regions in which toroidal fields dominate as well as regions in which poloidal fields dominate. In particular, the field seems to be toroidal in the denser molecular material that is concentrated near the Galactic plane, and poloidal in the more diffuse, hotter, and more tenuous synchrotron-emitting regions.
3 The model of Uchida, Shibata, and Sofue A theoretical model for the Galactic center that may be able account for the separate “poloidal-dominant” and “toroidal-dominant” regions that we see in Figure 2 is the magnetodynamic model developed by Uchida, Shibata & Sofue (1985), and further refined by Shibata & Uchida (1987). This model was developed in order to explain the “Galactic Center Lobe” (GCL), that is a limb-brightened radio structure with a size of several hundred pc extending from the plane of the Galaxy up towards positive Galactic latitudes (Sofue & Handa 1984). In the model of Uchida et al. (1985), the GCL represents a gas outflow that is magnetically driven. The model consists of nonsteady axisymmetric magneto-hydrodynamic simulations in which the field is assumed to be perpendicular to the Galactic plane at high Galactic latitudes, but acquires a toroidal component near the Galactic plane due to differential rotation of the gas to which it is coupled via flux-freezing. The stress of the resultant magnetic twist is what drives the outflow. In Novak et al. (2003), we argue that this model is fundamentally consistent with both the poloidal field seen by radio observers and the toroidal field seen by SPARO, provided that we make allowances for the clumpy distribution of the molecular gas.
4 Faraday rotation It is possible to probe the line-of-sight component of the magnetic field in any region of the Galactic center that contains thermal gas, provided that this region lies along the line-of-sight to a synchrotron source. This is because polarized radio emission suffers Faraday rotation as it passes through thermal gas. The model of Uchida et al. (1985) makes a specific prediction regarding the line-of-sight component of the magnetic field. To visualize this prediction, imagine dividing the Galactic center region into four quadrants
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A a (degrees) Fig. 1 Results of450 /mi polarimetry (bars) and photometry (contours) of the Galactic center. obtained using SPARO. The distribution of 450 pm flux closely follows the Galactic plane. that lies at a position angle of +31”. Coordinate offsets are measured with respect to the location of Sgr A* (that lies at the inmsection of the horizontal and vertical dotted lines). Each bar is drawn parallel to the inf‘eired magnetic field direction (i.e. perpendicular to the E-vector of the measured submillimeter polarization), and the length of the bar indicates the measured degree of polarization (sce key at bottom left). Contours are drawn at 0.075, 0.15, 0.30, 0.60. and 0.95 times the peak Rux, which is located at the position of Sgr 6 2 . For clarity, negative contours are not shown. The reference beam offsets were the same for polanmetry and photometry and are given in 3 2. The S’ beam of SPARO is shown in the key. Positive Galactic latitudes lie towards the upper right of the figure. and positive Galactic longitudes lie towards upper left.
according to the signs of Galactic longitude and latitude. According to the model, the sign of the line-ofsight field (i.e., towards or away from the observer), should be the same within a quadrant, and opposite in adjacent quadrants. We carried out a survey of the literature on Faraday rotation measurements toward Galactic center synchrotron soiirces, and we discovered a pattern of observed reversals in the sign of the Faraday “rotation measure” (hereafter, RM) that matches these predictions. We show this, using plus and minus signs, in Figure 3. The plus sign represents positive RM, and the minus sign represents negative RM. References to the Faraday rotation observations are given in Novak et al. (2003). We note that for RM measurements at positive Galactic longitudes, the asymmetry with respect to the Galactic plane had been noted previously, and Uchida et al. ( 1 985) pointed out the agreement with their model. To our knowledge, Novak et al. (2003)
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Act (degrees) Fig. 2 450 pm polarization measurements (bars) shown together with 90 cm radio continuum image (gray scale, LaRosa et al. 2000), and 850 pm continuum emission (contours, Pierce-Price et al. 2000). As in Fig. 1, the orientation of each bar is parallel to the inferred magnetic field direction (i.e., orthogonal to the measured direction of polarization) and its length is proportional to the degree of polariration. The radio continuum image shows about six locations where non-thermal filaments can be seen. These non-thermal filaments trace magnetic fields in hot ionized regions. The gray scale image is loganthmically scaled, and the contours of 850 pm emission are also logarithmic. Coordinate offsets are measured with respect to the position of Sgr A*. The location of the brightest bundle of non-thermal filaments (referred to as the non-thermal filaments of the Radio Arc) is indicated in the figure.
were the first to compare the signs of the RM measurements that lie towards negative Galactic longitudes with the predictions of the Uchida et al. ( I 985) model. The pattern that we see in Figure 3 is the one that results when the poloidal field points toward positive Galactic latitude. If the poloidal field is Laken to point in the negative latitude direction, then all of the RM signs should be reversed. Thus, if our interpretation of these data in the context of the Uchida et al. (1985) model is correct, we can conclude that the large scale poloidal field points towards Galactic North. Considering this pattern of RM reversals together with the SPAR0 polarization map and the evidence for poloidal fields derived from radio synchrotron maps, we conclude that the available observations do
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Fig. 3 Contours show distribution of 3.5 cm radio continuum emission from the central 500 pc (Haynes et al. 1992). The plus and minus symbols refer to the sign of the Faraday rotation measure, as discussed in 5 4.
support the general picture given by Uchida et al. (1 985) for the large-scale magnetic field in the Galactic center. The data, however, are very sparse. More observations, especially of Faraday rotation, are needed. The SPAR0 project was funded by the Center for Astrophysical Research in Antarctica (an NSF Science and Technology Center; OPP-8920223), by an NSF CAREER Award to G.N. (OPP-9618319), and by a NASA GSRP award to D.C. (NGT5-88). References Haynes, R. F., Stewart, R. T., Grey, A. D., Reich. W., Reich, P., & Mebold, U. 1992, A. & A,, 264, 500 LaRosa, T. N., Kassim, N. E., Lazio, T. J. W., & Hyman, S. D. 2000, AJ, 119,207 Morris, M., Davidson, J. A., Werner, M., Dotson, J., Figer, D. F.. Hildebrand, R. H., Novak, G.. & Platt, S. 1992, ApJ, 399, L63 Novak, G., Chuss, D. T., Renbarger, T., Griffin, G. S., Newcomb, M. G., Peterson, J. B., Loewenstein. R. F., Pernic, D., & Dotson, J. L. 2003, ApJ, 583. L83 Novak, G., Dotson, J . L., Dowell, C. D., Hildebrand, R. H., Renbarger, T., & Schleuning, D. A. 2000, ApJ, 529, 24 1 Peterson, J. B., Griffin, G. S., Newcomb, M. G.. Alvarez, D. L., Cantalupo, C. M., Morgan, D., Miller, K. W., Ganga, K., Pernic, D., & Thoma, M. 2000, ApJ, 532, L83 Pierce-Price, D., Richer, J. S., Greaves, 1. S., Holland, W. S., Jenness, T., Lasenby, A. N., White, G. J., Matthews, H. E., Ward-Thompson, D., Dent, W. R. E, Zylka, R., Mezger, P., Hasegawa, T., Oka, T., Omont, A,, & Gilmore, G. 2000, ApJ, 545, L121 Renbarger, T., Chuss, D., Dotson, .I.L., Hanna, J. L., Novak, G., Malhotra, P., Marshall, J., Loewenstein, R. F., & Pernic, R. 2003, in preparation Shibata, K., & Uchida, Y. 1987, PASJ, 39, 559 & Handa, T. 1984, Nature, 310, 568 Sofue, Y., Uchida, Y., Shibata, K., & Sofue, Y. 1985, Nature, 317, 699
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Astron. Nachr./AN 324, No. S 1, 139 - 143 (2003) / DO1 IO.IO02/asna.200385087
Extended photoionization and photodissociation in Sgr B2 J.R. Goicoechea*l,N.J. Rodriguez-Fernandez* * 2 , and J. Cernicharo***I
’ Departamento de Astrofisica Molecular e Infrdrroja, IEM/CSIC, Serrano 121, E-28006 Madrid, Spain
* Observatoirede Pans - LERMA. 61, Av. de I’Observatoire,75014 Paris, France Key words Galaxy: center - infrared: ISM: lines and bands ISM: individual (Sagittarius B2)
- ISM:
dust, extinction-H
II
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Abstract. We present large scale 9’ x 27’ (25 pc x 70 pc) far-IR observations of the Sgr B2 complex using the spectrometers on board the fnfrcired Spuce Obsen.atnr$ (ISO). The Par-IR spectra are dominated by the strong continuum emission of dust and by the fine structure lines of high excitation potential ions (N I I , N 111 and 0 111 ) and those of neutral or weakly ionized atoms (0I and C I I ). The line emission has revealed a very extended component of ionized gas. The study of the N I I I 57 p n / N I I 122 p,m and 0 111 52pm /88pm line intensity ratios show that the ionized gas has a density of 11,-.10~-~cm-3 while the ionizing radiation can be characterized by a diluted but hard continuum, with effective temperatures of ~ 3 5 , 0 0 0K. Photoionization models show that the total number of Lyman photons needed to explain such an extended component is approximately equal to that of the H 11 regions in Sgr B2(N) and (M) condensations. We propose that the inhomogeneous and clumpy structure of the cloud allows the radiation to reach large distances through the envelope. Therefore, photodissociation regions (PDRs) can be numerous at the interface of the ionized and the neutral gas. The analysis of the 0 I (63 and 145 pm) and C II (158 pm) lines indicates an incident far-UV field (Go, in units of the local interstellar radiation field) of 103-4 and a H density of cm-8 in such PDRs. We conclude that extended photoionization and photodissociation are also taking place in Sgr B2 in addition to more established phenomena such as widespread low-velocity shocks. ’
1 Introduction The Sgr B2 complex represents an interesting burst of massive star formation in the inner 400 pc of the Galaxy (the Galactic Center region, or GC) and may be representative of other active nuclei. Large scale continuum emission studies show that Sgr B2 is the brightest emission and the most massive cloud of the region (-.lo7 Ma;Lis & Goldsmith 1989). The main signposts of star activity are located within three dust condensations labelled as Sgr B2(N), (M) and (S). They contain all the tracers of on-going star formation: ultracompact H I I regions driven by the UV field of newly born OB stars, hot cores from embedded protostars, molecular masers, and high far-IR continuum intensity. These core regions are surrounded by a low density ( n ~ , < I 0cm-:’) ~ extended envelope (-lS’), hereafter Sgr 8 2 envelope, o f warm gas (Tk.200 K) and cool dust (Td=20-30 K; Hiittemeister et al. 1995). A summary of the different components present in Sgr B2 and their main charactcristics i s shown in Figure 1 (I@). The origin of the observed rich chemistry in the Sgr B2 envelope and its possible heating mechanisms are far from settled and several scenarios have been proposed. Low-velocity shocks have been traditionally invoked to explain the enhanced abundances of SiO or NH3 and the differences between gas and dust temperatures (cf. Floweret al. 1995). The origin of shocks in the Sgr B2 envelope have been associated * Corresponding author: e-mail: javierQdamir.iem,csic.es, Phone: +0034 91 561 68 00 ,Fax: +0034Yl 564 55 57 ** e-mail: Nemesio.Rodriguez-FernandezQobspm.fr * * * e-mail: cerniOdamir.iem.csic.es
@ LOO? WLEY-VCH Vzrlag GmhH b Co KGaA. Wriohcim
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Fig. 1 Rzght : Large scale IRAS image at 60 p n (Gordon et al. 1993) and IS0 target positions across Sgr B2 region. L e f t . Sketch showing the different structures and components in the Sgr B2 complex Hot cores are shown black shaded and H I I regions are the structures enclosing the stars. (Adapted from Huttemeister et al. 1995)
either with large scale cloud-cloud collisions or with small scale wind-blown bubbles produced by evolve massive stars (Martin-Pintado et al. 1999). The effect of the UV radiation in the Sgr B2 envelope has been traditionally ruled out because of the gas and dust temperature differences, the unusual chemistry and the abscense of thermal radio-continuum and ionized gas outside the H 11 regions and hot cores within the central condensations. Our observations reveal the presence of an extended component of ionized gas detected by fine structure emission. All this new data suggest that UV radiative-type processes are also important in the heating of the Sgr B envelope in addition to mechanical mechanisms, as can be in other GC clouds (see Rodriguez-Fernandez et al. 2003). In this contribution we present a brief summary of the results obtained by ISO' in the Sgr 8 2 envelope (Fig. 1 [right])concerning the ionized gas and the effects of the UV radiation.
2 Extinction corrections The large HZ column density (up to loz5cmP2) found in Sgr B2 suggests that even in the far-IR, fine structure lines can suffer appreciable reddening. We have estimated the prevailing extinction in each position by converting the continuum opacity into visual extinction. The spectra cannot be fitted with a single gray body. Thus, we have modeled the observed continuum spectrum as a sum of two gray bodies. The total continuum flux in the model is: S A = (I - e-TTorm) L3A(Twarm) nu,,,
+ (1
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BA(Tcold)flcold
(1)
where Bx(T,) is the Planck function, 7: is the continuum opacity and 0, is the solid angle subtended by the i dust component. The continuum opacity is given by r~ = 0.014 A" (30/A)' where /3 is the grain emissivity exponent of each dust component. The observed spectral energy distributions are best fitted with a dust component with a temperature of 13-22 K and a warmer component with a temperature of 24-39 K. The warmer component contributes less than 10 % of the total optical depth. The higher dust temperatures are those measured in the southern part of Sgr B2. Depending of the position, the derived extinction varies from -20 to -1000 magnitudes (see Table 1 for lower and upper limits to the visual extinction).
' Based
on observations with ISO, an ESA project with instruments funded by ESA Member States (especially the PI countries: France, Germany, the Netherlands and the United Kingdom) and with participation of ISAS and NASA.
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3 The ionized gas The far-IR spectra exhibit several fine structure lines from the ionized material. We have clearly detected the 0 I 63 and 145 pm, and C I I 158 pm lines in all observed positions. In addition, lines coming from the N 11 122 ILm, N I I I 57 l m , 0 111 52 and 88 pm lines are also detected in most positions, revealing a prominent component of ionized gas in the southern and eastern regions of Sgr B2. Table 1 lists the extinction corrected 0 I[I 52 pm/88 pm line intensity ratios for the two derived limits to the extinction across the region. For extended emission sources and lines excited by collisions with electrons (see Rubin et al. 1994) we derive electron densities (n,!) in the range ~ 1 0 ' ~c n" p 3 for the extended envelope. At the limited spectral resolution of the LWS/grating mode, 0 111 lines are hardly detected in the central positions. Nevertheless, Fig. 2b shows their Fabry-Perot detection in Sgr B2(M). Both 0 111 lines appear centered at V ~ s ~ z 5 15 0 km f s-' and do not show emission/absorption at more negative velocities (foreground gas). Considering A" > 1000 magnitudes, we found n,> lo"." cmp3 in Sgr B2(M). Table 1 also lists the extinction corrected N 111 57/N I[ 122 line intensity ratios. The minimum averaged ratio is 0.77 while the upper limits are dependent to the maximum extinction affecting the lines. For those ratios, we derive minimum effective temperatures (T,ff) for the ionizing radiation of 32.000 36.000 K . Those T,ff should be considered as a lower limit to the actual T,ff of the ionizing source if this is located far from the nebulargas. We have carried out CLOUDY (Ferland 1996) simulations showing that the observed line ratios are consistent with an scenario where almost all ionizing photons arise from the H 11 regions within Sgr B2(M) and (N). The total flux of Lyman continuum photons is 5-l and T,ff = 35,000 K . The differences in the observed N 111 /N 11 ratios are due to the dilution of the incident radiation (lower ionization parameter). Hence, the size of the ionized region can only be explained if the medium is highly inhomogeneous. This suggests that the clumpy nature of the cloud allows the radiation to reach large distances through the envelope. In this scenario, several PDRs can be expected in the interface between the ionized and the neutral gas. N
Table 1 Selected line ratios after correcting for the estimated minimum and maximum extiction. The different beam sizes of each LWS detector are taken into account and extended emission is considered. Offsets are in arsec. map
(0,-90) (0,-180) (0,-270) (0,-450) (0,-630) (0,-810) (270,O) (180,O) (90,O) (-90,O) (-I 80,O) (-270.0)
warm dust A V (mag)
1.2-1.5 1.6-2.0 1.1-1.8 5.2-8.0 0.7- 1.7 16-18 3.7-5.2 2.8-3.8 2.7-3.2 3.5-3.9 1.1-1.3 27-28 4.9- 6.6 7.4-9.0 28-34 21-26 2 1-25
cold dust A v (mag) 15-55 23-84 25-92 41-112 131-294 367-877 148-493 59-205 28-102 23-85 16-59 156-536 78-276 168-565 228-579 62- I68 45-124
0 111 R(52/88) 10 kpc, l3 7pG (Beck et al. 1994), NGC 4631 - z > 8 kpc, l3 = 2pG (Golla & Hummel 1994), NGC 891 and NGC 4561 - z 3 kpc, B = 1pG (Sukumar & Allen 1991). We will consider here the possibility that the observed excess is due to cosmic ray protons and study their propagation from some point sources located in the direction of the G C see also Bednarek, Giller & Zieliliska 2002). In particular, we have studied the point source image as a function of time. We have also studied its dependence on the model of the regular magnetic field in the Galaxy as well as on the proton energy. We have also considered whether the observed excess near the G C could be caused by particle emission from a single pulsar. N
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2 Propagation of EHE protons Protons ejected from a point-like source propagate in the Galactic magnetic field which consists of regular (Breg)and irregular (Birr)components. We have adopted two different models for the regular Galactic magnetic field: model I is that proposed and described in detail by Urbanik, Elstner & Beck (1997), model I1 is the bisymmetric field model with field reversals and odd panty (BSS-A) proposed by Han * Corresponding author: e-mail:
[email protected] @ 2003 WILEY-VCH Verlag GmhH B Co KGaA. Wcinhem
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Fig. 1 Amval directions of protons with energies 3 x 10'' eV injected by a point source in the GC (marked by the large dot) for model I (left) and model I1 (right). Maps (in galactic coordinates) from a) to f ) show directions of particles arriving in consecutive time delay intervals of 5 x lo3 yr, i.e. a) is for 0 - 5 x lo3 yr, _._,f) (2.5- 3) x lo4 yr, and from i) to 0) with intervals of 2 x lo4 yr, i.e. i) is for 0 - 2 x lo4 yr, ... , 0) (1 - 1.2) x lo5 yr. g) and o) Arrival directions integrated over time. h) and p) Delay time distribution of arriving particles; time in units of lo3 yr.
& Qiao (1994, see also Stanev 1997). We calculate numerically the proton trajectories within the range of energies corresponding to those of the AGASA excess eV) and the SUGAR excess - 101s.5 eV) from the direction of the GC. For 1.2 x lo6 protons ejected isotropically from a point-like source located at the Galactic Center we record the parameters (numbers, directions, and arrival times) of particles intersecting a sphere with the radius of 250 pc centred on the Earth. These events are considered as observed by a detector on the Earth. The numbers of the arriving particles as a function of travel time (the time of flight along the straight line being subtracted) are displayed in the form of histograms in Figs. Lh, 2h, and 3h for model I, and in Figs. lp, 2p, and 3p for model 11. It becomes evident that the distribution of the arrival times of particles with energies by a factor 2-3 larger is completely different. The bulk of particles arrive to the observer within 2.5 x lo4 years (model I) and 5 x lo4 years (model 11) for 3 x 10" eV (Figs. Ih and lp), up to ,-- lo5 years (model I) and lo6 years (model 11) for lo1* eV (Figs. 3h and 3p). In Figs. 1,2, and 3 from a) to 9 (model I) and from i) to n) (model 11) we show maps (in galactic coordinates, with longitude increasing to the left) with the arrival directions of protons intercepting the sphere around the Earth within consecutive time delay intervals chosen accordingly to the particle energy and magnetic field model (see figure captions). Maps summed up over time are in Figs. g) and o) showing the direction distribution in the case of a steady source.
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Let's first concentrate on the results for model I. The most interesting feature for protons with energies (2 - 3 ) x 10" eV is the particle clustering in multiple images of the source. These images appear at different places on the sky at different times after injection. A large number of protons reach the Earth's vicinity from directions close to the GC (shifted by about N loo towards positive longitudes) creating an
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eV. Maps for model I are for time delay intervals 2 x lo4 Fig. 2 As in Fig. 1 but for protons with energies 2 x yr, for model I1 - lo5 yr. Maps for model I1 start from l o 5 yr. extended source with the radius of about 20" -30". For protons with energies 10'' eV the amval directions become much more scattered (see Figs. 3a to 3h). In this case, protons arrive from a large part of the sky, almost independently of time after injection, apart from the peak for the first 2 x lo4 years. Increasing slightly the field strength will cause the discovered features shifting to higher energies and fitting better to the energy range where the actual excess of particles has been detected. The arrival directions of protons and their time distributions are completely different for the field model 11. There are no protons arriving directly from the actual position of the source at the GC (out of lo6 ejected) up to 3 x 10'' eV. For protons with 3 x eV, only a single image of the source is visible at a high negative latitude. It is created mainly by protons amving with relatively small time delay with respect to the rectilinear propagation, i.e. within less than 2 x lo4 years (see Fig. lp). For lower proton energies the image of the source is also centred on high galactic latitudes becoming broader and stronger. Particles arrive to the Earth much later than for model I i.e. after (2 - 4) x lo5 years for 2 x eV and ( 2 - 7) x lo5 years for 3 x 1018 eV. In spite of the instantaneous injection the anisotropy due to the source would be visible in the same directions for a long time. By comparing our calcufation results on the proton anisotropy for the two magnetic field models we conclude that the propagation of charged particles is very sensitive to their energies and to the structure and strength of the magnetic field. The two models, both based on experimental observations, give totally different predictions concerning the particle angular distribution on the sky, meaning that one should be very careful with drawing any conclusions based on one particular model of the Galactic magnetic field.
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3 Single pulsar as a plausible source We assume that particles with different energies are ejected isotropically by a short lived source, most likely a very young pulsar. Such very young pulsars (with milisecond periods) are presumably formed during the
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Fig. 3 The differential spectra of neutrons (dashed curves) and y-rays (full curves) observed from the Galactic Center on Earth at the time 10, lo2,lo3, 3 x lo3, and lo4 yrs (from the thinnest to the thickest curve) after the formation of the pulsar. The thick dotted curve shows schematically the observed cosmic ray spectrum within the 20n circle.
supernova type Ib/c explosions. The precursors of these supernova types are probably lower mass WolfRayet or oxygen-carbon stars rotating very fast and which have small mass explosion envelopes. Let's assume that at least one of them had parameters allowing acceleration of protons to energies above 10l8 eV. We follow the suggestion that pulsar winds are able to accelerate particles to energies E corresponding to the full potential drop available across the polar cap region (Gunn & Ostriker 1969, Blasi, Epstein & Olinto ZOOO),
where 0 = 27r/P, P = 10p3Pn,, s is the pulsar period, R = lo6 cm is the radius of the neutron star, B = 1013BB13 G is its surface magnetic field, e is the elementary charge, and c is the speed of light. The above equation allows us to constrain the parameters of the pulsar able to accelerate protons to energies E 2 10'' eV. The following condition has to be fulfilled P,, 5 8B:,/2.If the pulsar loses its rotational only on emission of the dipole radiation with the power L, then its period at specific time t is energy, ErOt, determined by the equation
Erot,= L=+IO(~
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where 1 = 1.4 x g cmP2 i s the neutron star moment of inertia. As a result of this energy losses the period of the pulsar changes in time according to P&(t) = Po",,,,, + 3.5 x 10p8B?,t, where f is in seconds, and f i ~ , ? is ~ the, pulsar initial period. By using the above equations, we estimate the time elapsed from the pulsar formation, t,,,, during which protons will be accelerated above 10Is eV (assuming that Po.,,, s,r//Bls 3 x lop3. We can not, however, exclude a possibility that the observed excess is being produced by even more delayed protons with a smaller intensity and a longer arrival time At. If the total observed CR flux was not produced by pulsars (but by some other sources) then such a hunch of particles could produce an increase on the average CR flux. For the bisymmetric field model proposed by Han & Qiao (1994) the image of the source located at the GC does not appear in the direction to the real source up to 3 x 10" eV, hut is shifted from its real position by a large angle. Therefore, in this case protons can not be responsible for the observed excess. N
4
Conclusions
1) Protons with energies between (1 3 ) x 10'' eV injected instantaneously by a point-like source at the Galactic Centre can form multiple images at directions completely different from those towards the source, as well as images shifted only slightly from the position towards the source. 2) The results of particle propagation for the two considered magnetic field models give totally different predictions for the particle angular distribution on the sky. ~
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3) The application of the Han & Qiao magnetic model produces strong north-south anisotropy of parti(2 - 3) x 10l8 eV. Most of the delayed particles arrive from the southern Galactic
cles with energies
N
hemisphere. 4) Our results do not depend on the particular distribution of the irregular magnetic component, providing that its magnitude is not larger than assumed here. 5) The maximum proton fluxes predicted for a pulsar model are larger than that observed. Thus, the model of isotropic and instantaneous particle injection by a pulsar, located close to the GC, could explain the observed flux of particles in the AGASA-SUGAR excess.
Acknowledgements The work is supported by the grants: KBN No. 5P03D02521 and the University of t 6 d i .
References Beck, R. et al., A&A 292,409 (1994). Bednarek, W., M. Giller, M. Zielinska, J.Phys.G, 28,2283 (2002). Bellido, J.A. et al.,Astropart.Phys.15, 167 (2001). Blasi, P., Epstein, R.I.,Olinto, A.V., ApJ 533, L123 (2000). Goldreich, P., Julian, W.H., ApJ, 157,869 (1969). Golla, G., Hummer, R., A&A 284,777 (1 994). Han, J.L., Qiao, G.J., A&A, 288,759 (1994). Hayashida, N., et al., AstropactPhys. 10,303 (1999). Stanev, T., ApJ 479,290 (1997). Sukumar, S., Allen, R.J., ApJ 382, 100 (1991). Urbanik, M., Elstner, D., Beck, R., A&A 326,465 (1997).
Astron. NachrJAN 324, No. S1. 151 - 155 (2003) / DO1 10.1002/asna.200385029
Discovery Of New SNR Candidates in the Galactic Center Region with ASCAand Chandra
',
Atsushi Senda* Hiroshi Murakami2,and Katsuji Koyama'
' Cosmic Ray Group, Dept. of Physics, Kyoto University, Sakyo-ku, Kyoto, 606-8502, Japan
' Institute of Space and Astronautical Science (ISAS),3-1 - 1 Yoshinodai, Sagamihara,Kanagawa 229-85 10, Japan
Key words X-ray, ISM, supernova remnamts, G0.570-0.018 PACS 04A25
We report the discovery of diffuse X-ray features which are possible SNR candidates near the Galactic Center (GC) observed with ASCA and Chandra. G0.570-0.018 has extremely small (20" diameter) shell-like morphology. Its X-ray spectrum exhibits strong Fe-K line emission and is well fitted by an NEI model with a temperature of about 6 keV. These characteristics suggest that G0.570-0.018 is a quite young (t 100 year) SNR. Diffuse hard X-rays were also detected from G359.92-0.09. Its X-ray spectrum also exhibits strong Fe-K line emission. The X-ray excess coincides with the shell-type feature observed in the radio continuum (e.g.. Ho et al. 1985), which is attributable to a new SNR. In addition, we have discovered several soft X-ray clumps. Their X-ray spectra are thermal (kT I keV) and clearly show atomic line features such as Si, S, Ar and Ca. The origin of the diffuse X-ray emission from the GC region has been an unresolved issue for over a decade. Hard clumps such as (30.570-0.018 are likely to be young/middle-aged SNRs, and could produce the hot component of the GC plasma, while relatively soft (- 1 keV) clumps, which also may be SNRs, could contribute to the cool component of the GC plasma.
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1 Introduction The origin of the diffuse X-ray emission from the Galactic center (GC)region has been an open issue for over a decade. Giiiga and ASCAfound a large-scale (1" x 1.8") thin-thermal plasma with strong line emissions from highly ionized atoms (Koyama et al. 1989, 1996; Yamauchi et al. 1990). The total thermal energy of the plasma is as large as ergs. Koyama et al. (1996) proposed that the plasma was created either by an energetic explosion that occurred at the central massive black hole (Sgr A * ) or by multiple 105 year. supernova explosions that took place within the past With its superior spatial resolution, C h a n d ~ ahas successfully resolved thousands of point sources in the GC region. However, most (- 90%) of the X-ray flux from the GC region is attributable to a diffuse component (Ebisawa et al. 2002; Wang et al. 2002). On the other hand, Chandvu also has revealed that the diffuse X-rays from the GC region are rather clumpy (Bamba et al. 2002). The presence of the clumpy structures may favor a multiple-SNe scenario. In fact, new X-ray supernova remnants (SNRs) have been discovered with Chandra (e.g. Sgr A East; Maeda et al. 2002). In this paper, we investigate newly discovered clumpy structures near the GC, which reveal the origin of the diffuse X-ray emission.
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' Corresponding author: e-mad: sendaQcr.scphys.kyoto-u.ac.jp,Phone. +8175 753 3869, Fax: +8175 753 3799 @ 2007 WILEY-VCH Verldg GmhH & Cu KGaA Weinhem
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2 Observations We used the archive data of seven fields of views (FOVs) of Chandra ACIS-I observations; Sgr B2, Sgr A*, and five FOVs of the Chundru GC Survey. The on-axis position and total exposure time of each observation are given in Tab. 1 . Table 1 Chandra ACIS observations
Target Name Sgr A* Sgr B2 GCS 13 GCS 14 GCS 16 GCS 17 GCS 19
Position (1, b) [deg] (359.94, -0.05) (0.59, -0.02) (0.00, -0.20) (0.00,0.00) (359.80, -0.20) (359.80,O.OO) (359.61, -0.20)
Exposure [sec] 48720 98989 10762 10762 10762 11261 11261
3 Results and Discussions 3.1 G0.570-0.018 G0.570-0.018 was discovered by ASCA (Sakano et al. 2002). Chandra observations resolved this source as a small shell-like structure with a diameter of -20”. The X-ray spectrum exhibits an extremely strong Fe-K line emission with an equivalent width of about 4 keV. The X-ray spectrum is well reproduced by a high temperature (- 6 keV) thin-thermal non-equilibrium ionization (NEI) model; hence the source is likely to be a young SNR. A detailed discussion of (30.570-0.0 I8 is given by Senda et al. (2002). 3.2 G359.92-0.09 Ho et al. (1985) detected a non-thermal radio continuum filament, called the “wisp,” 4’ south of Sgr A*. In addition, the inward curve of Sgr A East and the presence of other condensations (including the “wisp”) arranged in a circular shape imply the presence of a shell-like slructure, as shown by the solid circle in Fig. I. NH3 observations of the dynamics of nearby molecular material support this morphology, hence Coil et al. (2000) identified the source, G359.92-0.09, as a new SNR candidate. As shown in Fig. 1, the Chandra observations show that an X-ray excess fills the eastern half (EH) and southwest part of the radio shell of (3359.92-0.09 (Murakami et al. 2002). Although the northwest part shows no clear excess within the shell, this is due to the contamination of the intense X-ray emission from the SNR Sgr A East (Maeda et al. 2002). On the southwest edge of the shell, an X-ray bright filament is also discovered, which clearly corresponds with the non-thermal “wisp.” We have extracted X-ray spectra from three different regions; Eastern half (EH), Southwest quadrant, and the “wisp.” A thermal NEI model yields an acceptable fit for a spectrum from each region (Fig. 2 and Tab. 2). The observed emission measure from the EH is I . 1 x cmP3 using a distance of 0=8.5 kpc for G359.92-0.09). From the radius of the shell (2’ 5 pc), the total volume of the plasma is determined to be T/toral 1.6 x cm3. Assuming that the half of the X-ray emission comes from the EH, we calculate the electron density and the thermal energy of the EH to be n, 0.4 and EEH= 37&,kTVEH 1.6x 1050 ergs, where kT is the best fit value for the EH spectrum. Even though we cannot estimate the X-ray properties of the NW quadrant, a crude extrapolation of the EH result suggests that the thermal energy from the whole shell would be reasonably expected from a typical supernova explosion. Assuming that the expansion velocity of the SNR shock front is the sound velocity (1000 km s-l at 10 keV), the age of G359.92-0.09 is
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Fig. 2 The X-ray spectrum of the Eastern-Half (EH) of the Fig. 1 The C h a n d ~ aACIS grayscale image of (359.92-0 .09 in the 3.0-8.0 keV band. The overlaid contours are from the VLA,6cm continuum map (Ho et al. 1985).
G359.92-0.09 obtained with the Chandra ACIS. The best-lil spcctrum with the NEI model is shown by solid line.
Table 2 Best-tit Results of the X-ray Spectra o f G3.59.92-0.09 with Chamfra
kT (keV) T (10’’ s cmp3) NH cm-’) Fluxn (10-12 erg sP1 cm-2 )
Eastern Half 11.4 (>3.4) 1.4 (0.8-5.4) 6.0 (3.7-8.3) 1.5
Southwestern Quadrant 2.6 (> 1.9) 2.1 (>0.01) 18 (7.0-33) 0.4
”wisp” 12.7 0 7 . 9 ) 88 (>0.001) 37 ( 3 2 4 4 ) 0.4
An NEI model was employed to produce the estimated wlues. Quantities in parentheses are 90% confidence limits. “Flux (no correction for absorption) i n the 2.0-10.0 Lev hand.
determined to be 3800 year. On thc other hand, the ionization parameter indicates that the age of the plasma of the EH is 1.2 x I O4 year, although the uncertainties of this estimate is largc. N
3.3 G359.79-0.26 and (3359.77-0.09 The soft band image from the Charrdra GC Survey shows that diffuse emission extends from Sgr A East to the southward direction. The extended emission is relatively soft and clumpy (Fig. 3). From these clumpy structures, we havc identified three prominent soft clumps named G359.79-0.26, (3359.77-0.09, and (3359.73-0.35. The X-ray spectrum of each clump exhibits K-line emissions from He-like and/or H-like ions of Si, S, Ar, and Ca. These clumps were also detected previously by ASCd and ROS.4T. We have combined the X-ray spectra and tried to fit them with one model. The best fit results arc shown in Fig. 4 and Tab. 3. The large (NH 5 x lo2’ cm-2) absorption columns of (3359.79-0.26 and G359.77-0.09 indicate that they are located near the GC, while the significantly smaller ( N H lo2’ cm-’) absorption column of (3359.73-0.35 suggests that this clump is a foreground object. The results or spectral fitting with a thermal NEI model show that physical parameters of G359.79-0.26 and (3359.77-0.09 (NHand metal abundances) are similar to each other. In addition, the 1-3 keV band image suggests that (3359.79-0.26 and G359.77-0.09 are the southeastern and northwestern parts of a largc (- 30 pc) elliptical shell. This indicates that the two clumps have the same origin, an energetic N
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The image Of the ’Oft (1’0-3’0 kev) band near the GC region. The image center corresponds to an on-axis position of the FOV of GCS16.
Fig. 4 Combined X-ray spectrum of G359.79-0.26 obtained with ACIS-I (black), A S C A GIS2 (red), GIS3 (green), and ROSAT PSpC (blue), Solid lines show the best-fit thermal NEl model,
explosion such as a supernova, which occured at the center of the large shell. However, their temperatures are slightly different, so the interpretation of their origin is still preliminary. Table 3 Best-fit Results of the Combined X-ray Spectra of the Soft Clumps
kT (keV) T (10’~ s~ m - ~ ) N H (10” cm-’) FIUX~ ergs-’ cm-2 -Abundances (solar)-
Si S Arc Ca
)
G359.79-0.26 G359.77-0.09 (3359.73-0.35 0.84(0.75-0.93) 1.31 (1.03-1.79) 1.4(1.12-1.52) 0.5 (0.1-1.6) 9.9(-) 64 (>5.6) 4.9(4.65.2) 5.8 ( 5 . 1 4 . 5 ) 1.2(0.8-1.4) 2.3 2.2 I.Ib 0.4(0.3-0.6) 0.8(0.5-1.0) 1.3 (0.4-2.2) 3.2(1.4-5.6)
0.6(0.4-0.9) 0.8(0.5-2.1) 0.3 (< 1.0) 1.4 (< 3.3)
2.1 (1.2-3.8) 5.4(359.3) I 1 (3.1-27) 9.3(< 34)
Values in parentheses are 90% confidence limit\. aFlux (no correction of absorption) in the 2.&10.0 keV band. bFlux obtained with the ASCA GIS data. C A VMEKAL model is applied to obtain the results.
4 0
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Summary
Combining Chandra observations and archive data from ASCAand ROSAT,we have discovered several X-ray clumps in the GC region. Some of these clumps ((30.570-0.018 and G359.92-0.09) show thermal spectra from high temperature (-1 0 keV) plasmas; others ((3359.77-0.09 and G359.79-0.26) show thermal spectra from lower temperature (- 1 keV) plasmas. The X-ray emission from G359.92-0.09 is a counterpart to the non-thermal radio shell. Its energetics suggest that (3359.92-0.09 is a young/middle-aged (3800-1.2~lo4 year old) SNR.
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The X-ray properties of G359.77-0.09 and G359.79-0.26 appar to be typical of Galactic SNRs, but the natures of these two objects are still uncertain. The X-ray spectra of G0.570-0.018 and G359.92-0.09 are similar to that of the hard component of the GC plasma, while the spectra of (3359.77-0.09 and G359.79-0.26 are similar to that of the soft component. These characteristics suggest that the GC plasma may be largely resolve into individual clumps, which arc likely SNRs. However, the total diffuse emission from the GC region is greater than the sum of detected GC SNRs by 1-2 orders of magnitude.
References Bamba, A,, et al. 2001, PASJ 53, L21 Bamba, A,, et al. 2002, Proc. of “New Visions of the X-ray Universe in the X M M - Newton and Chandra era”, in press (astro-ph/0202010) Coil, A. L, & Ho, P. T. P. 2000, ApJ 533, 245 Ebisawa, K., et al. 2002,Proc. of “X-ray Surveys in the light of new observations”, in press (astro-ph/0210681) Koyama, K., et al. 1989, Nature 339, 603 Koyama, K., et al. 1996, PASJ 48,249 Ho, P. T. P., et al. 1985, ApJ 288, 575 Maeda, Y., et al. 2002, ApJ 570, 671 Murakami, H. 2002, Ph.D thesis, Kyoto University. Sakano, M., et al. 2002, ApJS 138, 19 Senda, A., Murakami, H., & Koyama. K. 2002. ApJ, 565, 1017 Wang, Q. D., et al. 2002, Nature 415, 148 Yamauchi, S., et al. 1990, ApJ, 1990, 365,532
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Astron. Nachr./AN 324. No. S1. 157- 160 (2003) / DO1 10.1002/asna.200385030
Molecular Line Observations of the Tornado Nebula and its Eye J. Lazendic’
‘.536,
M. Burton’, F. Yusef-Zadeh3,M. WardIe4, A. Green’, and J. Whiteoak‘
’ Harvard-SmithsonianCfA, 60 Garden Street, Cambridge MA 02138, USA
’ School of Physics, University of New South Wales, Sydney NSW 2052, Australia Department of Physics and Astronomy, Northwestern University, Evanston, IL 60208, USA Department of Physics, Macquarie University, NSW 2109, Australia School of Physics, University of Sydney, Sydney NSW 2006, Australia Australia Telescope National Facility
Key words shock waves, molecular clouds, supernova remnants, star formation, Tornado Nebula, G357.70.1, Eye of Tornado, G357.63-0.06 PACS 04A25 We present millimetre and NIR molecular-line observations of the Tornado Nebula and its Eye. The observations were motivated by the presence of OH( 1720 MHz) maser emission towards the nebula, believed to be an indicator of interaction between a supemova remnant and a molecular cloud. We found that the distribution of molecular gas around the Tornado complements its radio morphology, implying that the nebula’s appearance has been influenced by the structure of the surrounding molecular gas. Our NIR HZ ObSerVdtions revealed the presence of shocked molecular gas at the location where the nebula is expanding into the surrounding molecular cloud. It has been suggested that the Eye of the Tornado is related to the nebula on the basis of their apparent proximity. Our NIR and millimetre-line observations show that the two objects are not spatially related. Bry line emission, in conjunction with IR data at longer wavelengths and high-resolution radio continuum observations, suggests that the Eye is a massive protostellar source deeply embedded within a dense molecular core.
1 Introduction The Tornado nebula (G357.7-0.1) is a peculiar radio source located towards the Galactic Centre region. It has been classified as a supernova remnant (SNR) due to its steep radio spectrum and linear polarization (e.g., Kundu et al. 1974, Caswell et al. 1980, Shaver et al. 1985a), hut its unique morphology has led to other interpretations (e.g., an accretion powered nebula, Becker & Helfand 1985). The Eye of the Tornado (G357.63-0.06) is a compact radio source located 30” from the emission peak of the nebula. It was initially thought to he responsible for the formation of the nebula (e.g., through mass ejection from a pulsar or accreting binary system), hut was instead found to have a flat radio spectrum and was suggested to he an HII region (Shaver et at. 1985b).
2
Tornado
Frail et al. (1996) found a single OH(1720 MHz) maser at the northwestern tip of the Tornado (see Figure 1). When not accompanied by maser emission from the other three OH ground-state transitions at 1612, 1665 and 1667 MHz, the detection of this maser has been recognized as a signature of SNWmolecular * Corresponding author: e-mail: jlazendicQcfa.harvard.edu
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cloud interactions (see Koralesky et al. 1998 and references therein). Its presence may support the classification of the nebula as an SNR. The maser has a velocity of -12.4 km s-', implying a distance of 11.8 kpc to the nebula and placing the Tornado behind the Galactic Centre. 52'
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Fig. 1 (lefr) A 20 cm VLA image of the Tornado Nebula. The square marks the part of the nebula covered by the U N S W observations. frighr) Contours of' H2 line emission superimposed on a greyscale 20 cm radio continuum image of the northwestern part of the Tornado. The contour levels are: 1.6, 4.7, 6.3, 7.9, 9.5, 1 1 . 1 , 12.6, 14.2 x ergs-' cm-2 sr-'. The white and black crosses mark the location of the OH(1720 MHz) maser.
Using the University of New South Wales Fabry-Perot narrow-band tunable filter (UNSWIRF), we detected 2.12 pm Hz 1-0 S(l) emission towards the OH(1720 MHz) maser in the Tornado Nebula (see Figure 1). The correlation of the emission peaks in the radio continuum and H2 images suggest that the H2 emission originates from an expansion of a shock wave and is most probably shock excited, as found in other SNRs associated with the OH( 1720 MHz) maser (e.g., Lazendic et al. 2002a,b). The OH( 1720 MHz) maser is located at the western edge of the Hz emission, which is more sharply defined than the rest of the ring, probably delineating the leading edge of the shock front. Molecular transitions at millimetre wavelengths were also detected at the maser velocity of -12 km s-l using the 15-m Swedish-ESO Submillimeter Telescope (SEST). Emission from molecular species other than l 2 C 0 and 13C0, e.g., HCO+, HCN and H2C0, was found to be very weak (see Lazendic et al. 2003 for more details). Molecular gas associated with the OH(1720 MHz) maser and HP emission is optically ) . density is in agreement with the requirements for the thick, cold (-7 K) and dense (- lo5 ~ m - ~ This OH(1720 MHz) maser production in the post-shock gas behind the SNR shock front (Lockett et al. 1999), but the temperature is much lower than that expected in the post-shock gas in which the maser is created (50 - 125 K). However, since the cloud is optically thick, our CO observations are probing only the envelope of the cloud. Observations of more optically thin transitions of "CO and 13C0 are needed to examine the whole cloud temperature. The structure of the associated molecular gas complements the radio morphology of the Tornado Nebula (see Figure 2), implying that the distribution of the surrounding medium has influenced the nebula's unusual appearance. In particular, two minima in the molecular gas distribution, located symmetrically on each side of the nebula, coincide with large arc-like filaments in the nebula and point to locations where the shock could perhaps expand more readily than in the other regions.
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Fig. 3 NIR and radio images of the Eye. ( I t $ ) 2.16pm continuum image of the field centred on the Eye, overlaid with the Bry line image. Contours are at 3, 5. 7, 9. 12, 15 and 16 x lo-'' W m-2 arcsec-'. (righr) 6 cm VLA image Jy beam-'. overlaid with contours of 20 cm VLA image. Contours are at 8,26, 52, 104, 156,208 and 258 x
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Using the 3.8-m UK Infrared Telescope in conjunction with the CGS4 spectrometer we found 2.16pm Bry emission towards the Eye. The emission peaks at a velocity of N -200 kms-l, which drastically differs from the velocity of the molecular gas associated with the Tornado. The velocity of the Bry emission indicates the distance to the Eye is 8.5 kpc, which makes the Eye foreground to the Tornado Nebula. A similar velocity towards the Eye has also been measured using the H92a radio recombination line (Brogan & Goss 2003). The Eye is resolved by the NIR and radio measurements (see Figure 3) as a compact HI1 region, and therefore must be undergoing massive star formation. It consists of four knots of emission, each about 1.5” across and of similar brightness, placed symmetrically about the perimeter of a 6” diameter circle. There are faint extensions -2” to the south and to the west in the Bry image, but no emission from the centre. A fit to flux measurements from our M R and other IR data obtained from Midcourse Space Explorer (MSX)and Znfrured Astronomy Satellite (IRAS)is consistent with the Eye being a warm (- 190 K). unresolved ( w 0.05”) blackbody source at the core of an extended ( w 5.5”), cold (- 35K) greybody. The best fit value of this two-component greybody gives an angular size for the Eye which is similar to the size derived from the NIR and radio images. The Eye’s integrated infrared luminosity of 2 x 104La suggests it harbors a massive (-12Ma) protostellar source, perhaps a BO star (see Burton et al. 2003 for more details).
-
References Becker, R. H. & Helfand, D. J. 1985, Nature, 313, 1 15 Brogan, C. & Goss, W. M. 2003, AJ, 125,272 Burton, M. G., Lazendic, J. S., Yusef-Zadeh, F. & Wardle, M. 2003, in preparation Caswell, J. L.; Haynes, R. F.; Milne, D. K.; Wellington, K. J. 1980, MNRAS, 190,881 Frail, D. A,, Goss, W. M., Reynoso, E. M., Giacani, E. B., Green, A. J. & Otrupcck, R. 1996, AJ, 11 I , 165 1 Koralesky, B., Frail, D. A., Goss, W. M., Claussen, M. J. & Green, 1998, AJ, 116, 1323 Kundu, M. R., Velusamy, T. & Hardee, P. E. 1974, AJ, 79, 132 Lazendic, J. S., Wardle, M., Burton, M. G., Yusef-Zadeh, F., Whiteoak, J. B., Green, A. J. & Ashley, M. C. B. 2002a, MNRAS, 331,537 Lazendic, J. S., Wardle, M., Green, A. J., Whiteoak, J. B. & Burton, M. G. 2002h, in “Neutron Stars in Supernova Remnants”, Eds. P. 0. Slam and B. M. Gaensler, p.339 Lazendic, J. S., Wardle, M., Burton, M. G., Yusef-Zadeh, F., Whiteoak, J. B., Green, A. J. 2003, MNRAS, in preparation Lockett, P., Gauthier, E. & Elitzur, M. 1999, ApJ, 51 1,235 Shaver, P. A., Salter, C. J., Patnaik, A. R., van Gorkom, J. H. & Hunt, G. C. 198Sa, Nature, 3 13, 113 Shaver, P. A,, Pottasch, S. R., Salter, C. J., Patnaik, A. R., van Gorkom, J. H. & Hunt, G. C. 1985b, A&A, 147, L23
A a o n Nachr./AN 324, No. Sl. 161 - 165 (2003)/ DO1 10 1002/aana.200385031
The Search for Water and Other Molecules in the Galactic Centre with the Odin Satellite Aa. Sandqvist*I , P. Bergman’, A. Hjalmarson’, E. Falgarone3,T. Liljestrom4, M. Lindqvist2, A. Winnberg?, and the Odin Team’ Stockholm Observatory, SCFAB-AlbaNova,SE-106 91 Stockholm, Sweden Onsala Space Observatory, SE-439 92 Onsala, Sweden Ecole NomialcSupCrieure,FR-75005 Pans, France Metsahovi Radio Observatory, Helsinki University of Technology,FIN-02 150 Espoo, Finland
’ ’ ’ ‘ ’ http://www.snsb.se/Odin/Odin.html
Key words HzO, Sgr A Complex, CND, +20 km s-l cloud, +SO km s-’ cloud PACS 04A25
Observations with the Odin midsubmni space telescope have been made towards the Sgr A Complcx (the CircumNLiclear Disk, the +20 and +SO kin s- molecular clouds) in the Galactic Centre and we report here on the results of searches for HaO, H2”O and other molecules in these regions.
1 The Odin Satellite Odin is a millimetreisubmillimetre wave spectroscopy astronomy and aeronomy satellite, launched with a START-I rocket on 20 February 2001 from Svobodny, Russia in far-eastern Siberia. It has a I .I-m highprecision telescope with a beam efficiency ol‘ about 90% and beamwidths ol‘ 2‘.1 and 9 ’ 5 at submm and mm wavelengths, respectively. Its pointing uncertainty is < 10“. The submm radiometer consists of four cryo-cooled submm receivers tunable in the frequency range of 486 - 580 GHz with a single sideband tcmperature of M 3000 K. A cryo-cooled HEMT receiver, which is tuned to 1 19 GHz and dedicated to the search for 0 2 , has a single sideband temperature of M 600 K. The backend spectrometers consist oi an acousto-optical spectrometer (AOS) with a total bandwidth of 1040 MHz and two auto-correlators with bandwiths in the range of 100 - 800 MHz, corresponding to velocity resolutions of 0.08 - 1 .0 km s p L . The satellite is described in detail by Frisk et al. (2003) and the receiver calibration by Olberg ct al. (2003). Figure 1 is a photo montage of Odin in Earth orbit. Odin is a Swedish-led satellile project funded jointly by the Swedish National Space Board (SNSB), the Canadian Space Agency (CSA), the National Technology Agency of Finland (Tekes) and the Centre National d’Etudes Spatiales (CNES, France). The Swedish Space Corporation was the prime contractor i’or development and launch of Odjn and is responsible for the operation of the satellite.
2 Odin Observations of the Galactic Centre The molecular complex associated with Sgr A consists predominantly of a molecular belt comprising the “+50 km cloud” (M-0.02-0.07), the “+20 km spl cloud’ (M-O.13-0.08), and the Circuinnuclear Disk *
Corresponding author: e-mail:
[email protected]: +46 - 8 5537 851 I . Fax: +46 - 8 5537 8510
@ 1001 WILEY-VCH M r l a g GmhH & Co. KGaA. Wrmhcua
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Fig. 1 The Odin millimetre/submillimetresatellite
(CND) which surrounds Sgr A West and has a rotational velocity of the order of 100 km s-' in the same direction as the rotation ofthe Galaxy (see Fig. 2). These warm and high-density Galactic Centre molecular clouds are intimately entwined and interact with the continuum complex of Sgr A (e.g. Sandqvist 1989; Zylka et al. 1990) Three positions towards Sgr A have so far been observed with Odin, namely Sgr A* with the CND, the +20 km s-' molecular cloud, and the +50 km spl molecular cloud. The observed positions are marked on Fig. 2. Observations have been made in the spectral lines of 119-GHz 02,487-GHz 0 2 , 492-GHz C I , 548-GHz Ha80, 557-GHz Hi60, 572-GHz NH3, and 576-GHz ( J = 5 - 4) CO. However, only the data for H i 6 0 and Hi8O have been fully calibrated and reduced so far and they are presented in the next section. These results have recently been published as part of a special Odin Letters issue of Astronomy & Astrophysics (Sandqvist et al. 2003).
3 Water in the Sgr A Complex and the Expanding Molecular Ring Strong emission and absorption lines have been observed in the Ha60 line at all three Sgr A positions and they are presented in Fig. 3. However, no spectral line features can be detected in the H;80 line towards Sgr A* CND down to the rms noise limit of x 0.02 K. The line of sight towards the Sgr A Complex also crosses the massive x 180-pc Expanding Molecular Ring (EMR) surrounding the Galactic Centre and various spiral arm features further out in the Galaxy. A Gaussian analysis has been performed on the Sgr A' CND H!j60 profile in Fig. 3 using four absorption components and two emission components. The continuum emission was first subtracted out by fitting a linear baseline to the outermost channels on either side of the profile. The results were: Feuture Source = (Velocity (km spl), T i (K), Halfwidth (kms-l)) - I CND = (f73.2, f0.32, 88.5); IZ CND = (-31.6,
Astron. Nachr./AN 324, No. SI (2003)
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RIGHT ASCENSION Fig. 2 Observed positions (epoch 1950.0) in the Galactic Centre Sgr A region are marked by circles whose diameter indicates the Odin beam (2'.1) - Sgr A* CND: (17"42"'29".3, -28"59'18"); +20 kms-' Cluud: (17"42"29*.3, -29"02'18"); +50 kms-' Cloud: (17h42"41s.0, -28"8'00''). The Complex consists of the 20-cm continuum radiograph, the CND (thin lines) in HCN, and the molecular belt in 2-mm HzCO (solid lines) with isovelocity contours (broken lines). The cross marks the position of Sgr A* (adapted from Sandqvist 1989).
f0.24, 47.9); III Local Sgr Arm = (-4.8, -0.24, 13.4); IV -30 kms-' Arm = (-30.2, -0.25, 11.0); V 3 - k p c A m = (-53.5, -0.21,8.1); VIEMI?= (-132.2, -0.09,60.0). The first two components, I and 11, both seen in emission, are believed to originate in the rapidly rotating CND. The northeastern part of the CND is receding and the southwestern part approaching, which gives the asymmetric, somewhat double-peaked line profile structure. The 2. I-arcmin beam of Odin encloses fully the CND and the resulting velocity structure of the profile is reminiscent of that seen in many other molecular lines. The two H 2 0 profiles towards the +20 and +50 km S K ' clouds in Fig. 3 are marked by the characteristic emission component from these molecular clouds at velocities near +20 km s-' and +50 km s-', respectively. Furthermore, a new molecular feature in the Galactic Centre can now been identified. It is detected as broad H i 6 0 absorption in the velocity range of M +120 to +220 kms-'. We shall call this feature the High Positive Velocity Gas (HPVG). This feature is not seen in the Sgr A* CND profile, which we interpret as being due to the background continuum emission seen at this position being somewhat lower than towards the dust continuum peak emission from the Sgr A +20 and +50 km sP1 molecular clouds (seen in 800 and 350 pm continuum maps of Lis & Carlstrom 1994 and Dowel1 et al. 1999, respectively). Evidence for the existence of the HPVG in the Sgr A region, seen in other spectral lines, is scarce although some is present in IS0 mid- and far-infrared observations of H 2 0 (Moneti et al. 2001), and VLA observations of H I (Dwarakanath et al. 2003, in preparation) and OH (Karlsson et al. 2003). The HPVG should not be confused with the molecular gas in the far side of the EMR whose velocity falls inside the
164
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Sandqvist et at.: HzO towards Sgr A
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Velocity [hm/s]
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Fig. 3 The Ha60 profiles towards the Sgr A Complex.
same range but whose emission lines are narrower. Also, the HPVG is seen in absorption which places it in front of the Galactic Centre continuum sources and thus it cannot be part of the far side of the EMR. The broad H2O absoption component (VI), seen at velocities near-132 km sP1, is observed in all three positions. This feature has its origin in the near side of the EMR.
4 Water Abundance in Spiral Arms Three narrow HzO absorption components, 111, IV and V, seen at velocities near -5, -30 and -53 kms-l, respectively, are observed at all three positions (see Fig. 3) and are well-known Galactic spiral arm features, which were first identified in early 21-cm H I observations. They originate along the line of sight crossing the so-called Local Sgr, -30 kms-' and 3-kpc spiral arm structures. The absorption feature (111) at -5 km s-l appears to be the strongest and, judged by estimated continuum levels, this feature has an optical depth of at least one. With this assumption, a lower HzO column density limit of > 2 x 1013 is found for feature 111 in the Sgr A* CND profile, using the appropriate line width and an excitation temperature of 15 K. We can estimate Hz column densities using the C180-profiles in Fig. 4 which are the C l 8 0 profiles resulting for the three H 2 0 positions from a convolution of the SEST Cl80 (1 - 0) survey of the Galactic Centre spectra (Lindqvist et al. 1995) to a resolution of 2' (corresponding to the Odin beam size).The integrated intensities have been determined over the regions corresponding to the three narrow H20 absorptions. The H2 column densities have then been calculated by assuming optically thin emission, an excitation temperature of 15 K and a C l 8 0 abundance of 2 x lop7 with respect to H a (Frerking et al. 1982). This results in an H2 column density of 1.0 x loz2 for feature I11 in the Sgr A* CND profile. The corresponding H2O abundance limit for this feature is then > 2 x loP9.
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?j:€P2 2 -
Our non-detection of Hg80 towards Sgr A* CND provides an upper limit on the HzO abundance in the narrow absorption features. Using a lsO/lsO ratio of 500 (Wilson & Rood 1994) for this local absorbing cloud (111) and an excitation temperature of 15 K, we obtain an upper limit for the H 2 0 column density of 5 x l O I 4 cm-'. Hence, for the Local Sgr Arm absorption, the 3a upper limit of the H 20 abundance becomes 5 x lo-* while the lower limit was found to be 2 x lo-'. The average H.20 abundance estimated for the foreground gas towards Sgr B2 by Neufeld et al. (2000) is 6 x 10W7, which is about an order of magnitude higher than our range towards Sgr A. On the other hand, our range is in better agreement with HzO abundances found in giant molecular cloud cores by Snell et al. (2000) and in a local diffuse molecular cloud by Neufeld et al. (2002).
References Dowel1 C.D., Lis D.C., Serabyn E., Gardner M., Kovacs A,, Yamashita S. 1999, in ASP Conf. Ser. 186, The Central Parsec of the Galaxy, eds. H. Falcke, W.J. Cotera, W.J. Duschl, F. Melia, M.J. Rieke, 453 Frerking M.A., Langer W.D., Wilson R.W. 1982, ApJ 262,590 Frisk U., Hagstriim M., Ala-Laurinaho J. et al. 2003, A&A, 402, L27 Karlsson R., Sjouwerman L.O., Sandqvist Aa., Whiteoak J.B. 2003, A&A 403, 101 1 Lindqvist M., Sandqvist Aa., Winnberg A., Johansson L.E.B., Nyman L-A. 1995, A&AS 113, 257 Linke R., Stark A.A., Frerking M.A. 1981, ApJ 243, 147 Lis D.C., Carlstrom J.E. 1994, ApJ 424, 189 Moneti A., Cernicharo J., Pardo J.R. 2001, ApJ 549, L203 Neufeld D.A., Ashhy M.L.N., Bergin G. et al. 2000, ApJ 539, L111 Neufeld D.A., Kaufman M.J., Goldsmith P.F., Hollenbach D.J., Plume R. 2002, ApJ 580, 278 Olherg M., Frisk U., Lecacheux A. et al. 2003, A&A, 402, L35 Sandqvist Aa. 1989, A&A 223,293 Sandqvist Aa., Bergman P., Black J. et al. 2003, A&A, 402, L63 Snell R.L., Howe J.E., Ashby M.L.N. et al. 2000, ApJ 539, L101 Wilson T.L., Rood R. 1994, ARA&A 32, 19 1 Zylka R., Mezger P.G. Wink J.E., 1990, A&A 234, 133
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Astron. Nachr.iAN324, No. Sl, 167- 172 (2003) / DO1 IO.lO02/asna.20038.51 10
Deep X-Ray Imaging of the Central 20 Parsecs of the Galaxy with Chandra
',
Mark Morris* Fred Baganoff', Michael Muno2, Christian Howard', Yoshitomo Maeda", Eric Feigelson4, Marshall Bautz2, Niel Brandt4, George Chartas4, Gordon Garmire4, and Lisa Townsley4
' Division of Astronomy, UCLA, Los Angeles, CA 90095-1562. USA Center for Space Research, Massachusetts Institute of Technology, Cambridge, MA 02139-4307, USA Institute of Space & Astronautical Science, 3- 1-1 Yoshinodai, Sagamihara, Kanagawa, 229-8510, Japan Department of Astronomy &Astrophysics, Penn. State Univ., 525 Davey Lab., University Park, PA 16802, USA
Key words Galaxy: center, X-rays: ISM, Sgr A*
Abstract. A deep observation toward the Galactic center with the Chandra X-Ray Observatory revealed a number of extended features, in addition to Sgr A* and SgrA East. Here, we focus on two curious, extended X-ray structures: large-scale (-10 pc) bipolar lobes centered on Sgr A* and a bright cometary source located 0.3 pc from Sgr A*, CXOGC 5174539.7-290020. The bipolar lobes consist of a number of emission clumps oriented along a line perpendicular to the Galactic plane, suggesting that a series of ejections has taken place on characteristic time scales o f hundreds to thousands of years. The clumps are embedded in a low-intensity, edge-brightened lobe which is most evidcnt in a flux ratio map. At two locations along the lobe, nonthermal linear features are present, suggesting that relativistic electrons may be impinging on the compressed, magnetic wall of this structure. The cometary X-ray source has no counterpart at other wavelengths; its orientation is consistent with a high-velocity neutron star ejected from the grouping of stars at IRSl3, but there are problems with that hypothesis, and other models warrant consideration.
1 Introduction The ACIS instrument aboard the Chandra X-Ray Observatory has been pointed at the Galactic center on several occasions since 1999, culminating in a deep observation in 2002 May-June. With a total useful exposure time of 590 ksec, a rich variety of phenomena are revealed, and are reported here: Sgr A*, its extent, its time variations and its possible jet (Baganoff et al. 2003a, 2003b), stellar point sources (Muno et al. 2003), SgrA East (Maeda et al. 2003), and extended continuum and line sources emission (Park et al. 2003). In this paper, we focus on a few extended X-ray sources of particular interest, both of which have been newly revealed by Chandra: bipolar lobes centered on Sgr A* and an unusual cometary source located near Sgr A* (Baganoff et al. 2 0 0 3 ~ ) .
2 The Bipolar Lobes Figure 1 shows a broad-band image of the center of the ACIS field of view, using data from 3.3 to 4.7 keV. The region around Sgr A*, including SgrA East, IS saturated in order to emphasize the fainter structures oriented perpendicular to the Galactic plane. A number of diffuse, extended structures are present in the image, but the most prominent of them are aligned along a line passing through Sgr A* and oriented perpendicular to the Galactic plane. Overall, these features appear to form a roughly point-symmetric * Corresponding author: e-mail:
[email protected],Phone: 1 310 825-3320, Fax: 1 3 10 206-2096
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structure centered on Sgr A*, reminiscent of bipolar nebulae around mass-losing stars. The bipolar lobes are apparently filled structures having in places a discernible sharp edge. The well-defined edge is better revealed by figure 2, which shows the flux ratio in two bands, emphasizing features with relatively soft spectra. The ratio map is affected somewhat by absorption by molecular clouds in the near foreground to the lobe emission (Howard, in preparation), which probably contributes at least partially to the asymmetry of the ratio map. Thermal fits to the ‘‘blobs’’ along the central axis of the bipolar lobes give temperatures of 2 keV, on average, with variations of a factor of 2 or so. The uncertainties in the lobe temperatures are considerable because of the difficulty of finding an appropriate local background to subtract. At present, no radial trends in gas temperature are evident.
Fig. 1 Broad-hand (3.3 - 4.7 keV) X-ray emission from a 35 x 28 pc region centered o n Sgr A * . The image has heen adaptively smoothed, and all point sources identified by Muno et al. (2003) have been removed. The true Galactic plane and the positioii of Sgr A * are indicated in black.
Along the northern edge of the northwestern lobe lie two linear X-ray features, shown in the unsmoothed image in Figure 3 as features “e” and “f”. These features have nonthermal spectra with power law indices of 1.3. Source “ d is very similar, but it bears no obvious relationship to one of the bipolar lobes. It is not clear whether these structures are truly filamentary, or whether they are bright by virtue of being edge-on surfaces. If the former view applies, thcn they are likely to be magnetic structures, by analogy with the nonthermal radio filaments located in this region. If the latter picture is applicable, then sources “e” and
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Fig. 2 "Softness Ratio" map of the bipolar lobes, consisting of the intensity ratio of the 2 - 4.7 keV band to the 4.7 - 8 keV band. The region shown is 25 x 44 pc (10 x 17.5 arcmin). For both bands, adaptively smoothed images are used, and all point sources meeting our selection criteria have been removed (see Muno et al. 2003). The dark spots represent unsnbtracted, hard sources. The central region around Sgr A* is not well represented by the point-source extraction routine because of the extremely high surface density of point sources there. Note the well-defined edge to both lobes.
"f" may be strong shocks occuring where relativistic particles occupying this lobe strike the compressed medium at the edge of the lobe. A compelling interpretation for the bipolar lobes straddling Sgr A* is that they result from energetic mass ejections from the immediate environment of Sgr A*, presumably from an accretion disk (see Baganoff et al. 2003a for a discussion of a possible jet which may be feeding the X-ray lobes). The 1 - 2 arcminute radial separations of the blobs corresponds to time intervals of (2500 - 5000 years)x(1000 km s - ' N ) , where V is the unknown outflow velocity of the X-ray emitting material. These characteristic times are somewhat less than the expected time between successive tidal disruptions of stars by the central black hole, 1-3 x lo4 years (Alexander & Hopman 2003), unless V is unreasonably small, so it may be difficult to ascribe the mass ejections to stellar tidal disruptions. The time scale for arrival at the outermost observable extent of the bipolar lobes is (1.5 x lo4) x (1000 km s-lN),which can be compared to the expansion time of SgrA East, -lo4 years (Maeda et al. 2003; Mezger et al. 1989; Uchida et al. 1998). This possible correspondence in time scale is interesting in the context of the suggestion of Maeda et al. (2002,2003) that the expansion of the shell of SgrA East past the location of Sgr A* might have provoked an accretion event in the recent past, although that event was hypothesized to have been only a few hundred years ago. The bipolar X-ray lobes have counterparts in the radio, but the detailed correspondence is better at long radio wavelengths (90 cm; Pedlar et al. 1989; Nord et al. 2003) than at 20 cm or shorter wavelengths (Yusef-Zadeh & Morris 1987). It is only with the most recent 90-cm data of Nord et a]. (2003) that the bipolar lobes have appeared to be anything else but an extension of the radio halo of the SgrA complex. The steep-spectrum, low-frequency radio emission can be interpreted as synchrotron emission from a population of energetic electrons that are closely related to those responsible for the presumably thermal emission
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Fig. 3 Linear X-ray features near Sgr A*, displayed in an unsmoothed image of SgrA. The features labelled “a” corresponds to the “Plume” discussed by Baganoff et al. (2003~)and by Maeda et al. (2003), while feature “b” is the jet described by Baganoff et al. (2003a). Feature c lies at the base of the nortbwest X-ray lobe, and is coincident with the radio “streamers” discussed by Yusef-Zadeh & Moms (1987). Features d, e, and fare all short, bright segments having nonthermal spectra with power law indices of 1.3.
in the X-ray lobes. The total amount of mass involved in the X-ray lobes is quite small; the column density of electrons required to give the observed intensities, assuming a thermal emission model, implies a space density of only about 1 cmp3 if the line-of-sight depth of the blobs is equal to their widths on the sky. This, in turn, implies that each blob has only about 1 Mo.
3 The Cometary Source, CXOGC 5174539.7-290020 The central 18” x 20” of the Chandra field (Figure 4) shows a concentration of at least 4 bright sources, one of which is Sgr A*. Approximately 8” (0.3 pc) to the northwest of Sgr A* lies a remarkable ridge
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of X-ray emission. The intensity of this structure i s continuous along this ridge, increasing monotonically from 2” W of Sgr A* to a bright tip at IS”W, 8.5’’ N. At this point, we do not rule out the possibility that it consists of an unresolved string of 3 or more point sources, but there is no hint of that in the intensity image, and there is no evidence that the spectrum is also not continuous or homogeneous. The spectrum is well fit by a foreground-absorbed power-law having no indication of spectral line emission. With an overall spectral index of 1.7, it stands out as the source with the hardest spectrum in the immediate vicinity of Sgr A*. The absorption-corrected luminosity of this source is 1 .I x 1 034ergs s p l , almost 5 times larger than that of Sgr A* . The cometary geometry of CXOGC 5174539.7-290020 suggests a rapidly moving source of energetic particles, The spectral variations, shown in figure 4, are consistent with this, in the sense that the spectrum is hardest at the bright tip of the structure, and declines into the “tail”. This suggests that the source of particles is at the tip, that those particles are deposited in the source’s wake, and that they experience a net loss of energy as they radiate and distance themselves from the source. The nature of the source remains a mystery. There is no counterpart in the radio (e.g., Zhao & Goss 1998, Yusef-Zadeh, Roberts & Biretta 1998),in the mid-infrared (Morris et al., unpublished Keck data: see http://irastro.jpl.nasa.gov/GalCen/galcen.html), and the Gemini near-infrared images (Rigaut et al. 2003) show no obvious source associated with either the tip or the tail of CXOGC 5174539.7-290020. One clue may be provided by its orientation. The tail orientation is inconsistent with an origin based on linear ejection from Sgr A*; however, the tail does point back to the X-ray/infrared/radio source 1RS13, which appears to be a cluster of young, high-mass, stellar objects, possibly the surviving core of a substantial cluster (Maillard et al. 2003). A high-velocity neutron star could have been produced in such a cluster, either as the result of a supernova explosion, or by being gravitationally scattered out of the evaporating, tidally disrupting cluster (Kim & Morris 2003). If we posit that the speed of the neutron star is 300 km s-l in the plane of the sky, then the time scale for it to displace itself from IRS13 by the observed amount is I400 years. A supernova of this age should still be recognizable, even in this complex region, so ejection in some relaxation process would therefore be the more likely origin. Indeed, X-ray nebulae associated with pulsars can have a cometary morphology. For example, the pulsar B 1957+20, which has an X-ray powerlaw photon index of 1.9 f 0.5, has a cometary tail about half as long as that of CXOGC J 174539.7-290020 (Stappers et al. 2003), although the luminosity ofB1957+20 is three orders of magnitude smaller than the galactic center object. Furthermore, the absence of a radio source coincident with 5174539.7-290020 is mysterious if it i s truly a young neutron star. Consequently, other models warrant consideration. We mention three others here, although none of them are very compelling at this stage: 1) that the ridge results from the emission from hot electrons in a unidirectional jet emanating from some source in IRS13, 2) that the ridge is the laterally moving shock site where a rapidly precessing jet originating at Sgr A * impacts some (currently unseen) ambient medium (this model is inconsistent with the straightness of the jet on the opposite side of Sgr A*; Baganoff et al. 2003a), and 3) that the ridge is an essentially edge-on shock front resulting from a high-velocity wind, perhaps from the mass-losing stars in the IRS I6 cluster, impacting an ambient medium. Clearly, additional work remains to be done before the nature of this interesting cometary source can be clarified.
Acknowledgements This work has been supponed by NASA grants NAS8-00128, NAS8-38252 and G02-3115B.
References Alexander, T. & Hopman, C . 2003, ApJL, suhrnitted Baganoff, F.K., et al. 2003a, these proceedings Baganoff, EK., et al. 2003b. these proceedings Baganoff, F.K., et al. 2003c, ApJ, 590, in press Kim, S.S. & Moms, M. 2003, these proceedings Maeda. Y., et al. 2002, ApJ, 570, 671 Maeda, Y., et al. 2003, these proceedings Maillard, J.-P. et al. 2003, these proceedings
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r =1.55 +/- a.15
Fig. 4 Unsmoothed X-ray image of an 18” x 20” region showing the several sources near Sgr A* (marked with an “x”). The individual pixels have widths of 0.5 arcsec. CXOGC 5174539.7-290020 is the vertical ridge in the upper middle of the figure. The spectra of the entire source, and of the upper and lower parts, as indicated, are shown at the left, along with their power-law indices, r. Mezger, P.G. et al. 1989, A&A, 209,337 Muno, M., et al. 2003, these proceedings Nord, M.E., et al. 2003, these proceedings Park, S., et al. 2003, these proceedings Pedlar, A., et al. 1989, ApJ, 342, 769 Rigaut, F., et al. 2003, these proceedings Stappers, B.W., Gaensler, B.M., Kaspi, V.M., van der Klis, M. & Lewin, W.H.G. 2003, Science, 299, 1372 Uchida, K.I., Moms, M., Serabyn, E., Fong, D. & Meseroll, T. 1998, in IAU Syinp. No. 184: The Central Regions of the Galaxy and Galaxies, ed: Y. Sofue, (Dordrecht: Kluwer), p. 317 Yusef-Zadeh, F. & Moms, M. 1987, ApJ, 320,545 Yusef-Zadeh, F., Roberts, D.A. & Biretta, J. 1998, ApJL, 499, L159 Zhao, J.-H. & Goss, W.M. 1998, ApJL, 499, L163
Astron. Nachr./AN 324, No. S1, 173- 179 (2003) / DO1 10.1002/asna.200385099
Mapping Magnetic Fields in the Cold Dust at the Galactic Center David T. Chuss*', Giles Novak', Jacqueline A. Davidson3, Jessie L. Dotson', C. Darren Dowell', Roger H. Hildebrand6, and John E. Vaillancourt'
' NASA Goddard Space Flight Center ' Northwestern University
' USRA
NASA Ames Research Center
Caltech
' University of Chicago
' University of Wisconsin Key words Magnetic Fields, Galactic Center, Submillimeter Polarimetry, Dust
Abstract. We report the measurement of 158 new 350 pm polarimetry vectors in the central 30 parsecs of the Galactic center. These data were obtained at the Caltech Submillimeter Observatory using Hertz. Morphologically, these results show a consistency with previously published far-infrared and submillimeter results. We find that the angle of the magnetic field inferred from these observations is related to the 350 pm flux as obtained by SHARCKSO in the following way. At low fluxes, the magnetic field angle is consistent with that of a poloidal field as seen in nearby features such as the Galactic Center Radio Arc and the Northern and Southern Threads. At high fluxes, the magnetic field is oriented parallel to the plane of the Galaxy. This relationship suggests a model in which an initially poloidal field is sheared out in dense regions which are dominated by gravity. If this model is correct, it implies a characteristic field strength for the region of3 mG.
1 Introduction Ever since the discovery of the Galactic center Radio Arc (Yusef-Zadeh, Morris, & Chance, 1984), there has been much interest in the structure and strength of magnetic fields in the Galactic center. Even so, relative to other physical processes and structures in the Galactic center, comparatively little is know about the magnetic fields in this region. The study of the structure of the field was initiated by the observation that the Radio Arc and a number of other NTFs in the Galactic center (see LaRosa et al. 2000 and references therein) are oriented in a direction perpendicular to the Galactic plane. Such observations led to the hypothesis that the Galactic center is threaded by a global poloidal field. More recent results have given clues that the magnetic field structure in the Galactic center is more complicated than that of a simple poloidal field. Preliminary far-infrared observations of the Circumnuclear Disk (Morris et nl. 1988) indicated that the field running through regions of dust in the Galactic center does not trace a uniformly poloidal field. More recently, Novak et al. (2003) have used submillimeter polarimetry to trace the structure of the field in the central 200 parsecs and have found that the field in the cold dust on these scales is predominantly parallel to the Galactic plane. To reconcile their findings with the observations of the NTFs, these authors invoke a model (Uchida, Shibata & Sofue 1985) in which an * e-mail: chussQstars.gsfc.nasa.gov, Phone: 301 286 1858
@ 2003 WlLEY-VCH Veriag GmhH & Co KGaA. Wcmhem
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initially poloidal field can be sheared into a toroidal field in regions of relative high density by differential rotation and infall. The study of the strength of the field has a similarly interesting history. An estimate of the strength of the field associated with the Radio Arc was obtained by Yusef-Zadeh and Morris (1987). They estimated a lower limit of the field strength of - 1 mG by noting that the filaments maintain their structure against the ram pressure from interacting molecular clouds. The presence of such a large field strength led to two possibilities. First, it is possible that the mG field is the characteristic strength of a global field. However, the magnetic pressure implied by such a large global field strength is substantial and would require a large force for containment. If, on the other hand, the magnetic field in the filaments is enhanced with respect to that in the surrounding regions, the filaments would tend to dissipate on short timescales. There have been many Zeeman measurements of the line-of-sight field strength that have given values of between 10 pG and a few mG for various regions of the Galactic center (Plante, Low, & Crutcher 1995; Killeen, Lo, and Crutcher 1992). In some cases, such measurements provide a lower limit for the field strength, but without knowledge of the strength of the field in the plane of the sky, they do not help differentiate between the two ideas mentioned above. Far-infrared and submillimeter polarimetry provides a method for measuring the structure of the magnetic field as projected onto the plane of the sky. Rotating asymmetric dust grains become partially aligned such as to emit radiation in the far-infrared and submillimeter that is polarized in a direction perpendicular to the aligning magnetic field. Because of the extended distribution of dust in the Galactic center (Pierce-Price et al. 2000), this technique is an excellent way to obtain information about the structure of the magnetic field in this region. Unfortunately, because of uncertainties in dust grain physics and in the magnetic field component parallel to the line of sight, it is difficult to obtain a direct estimate of the magnetic field strength at any one point. However, using a model-based approach, we are able to obtain an indirect estimate of the magnetic field strength by looking at the overall structure of the measured field. We present submillimeter polarimetry of the central 30 pc of the Galactic center. We find that the magnetic field direction is neither entirely poloidal nor toroidal, but rather, the direction of the field depends on the submillimeter flux in the following way: for low fluxes, the field has an orientation perpendicular to the Galactic plane (poloidal); for high fluxes, the field’s orientation is parallel to the plane (toroidal).
2 Observations The 158 new polarimetry vectors were obtained at the Caltech Submillimeter Observator in May 2001. They are shown along with the data from Novak et al. (2000) in region 111of Figure 1. These data were obtained with the University of Chicago polarimeter Hertz, a 350 pm, 32 element array polarimeter with a ANA = 0.1 and FWHM resolution of 20”. For details of the observing procedure and data reduction see Chuss et al. (2003) and references therein.
3 Discussion The magnetic field structure inferred from all of the Hertz data obtained in this region to date are shown plotted over 850 pm continuum emission (Pierce-Price et al. 2000) in region I11 of Figure 1. Also shown are 60 pm polarimetry data (region I) and 100 p m polarimetry data (region 11) (Dotson et al. 2000). The field in the central 30 pc shows neither a predominantly poloidal nor toroidal direction. It does, however, appear to have structure on scales significantly larger than that of the molecular clouds (typically of the order of 5-10 pc). Also, the submillimeter data appear to trace magnetic fields that are spatially consistent with those traced by far-infrared polarimetry. Together, these observations indicate that there is some decoupling between the magnetic fields and the dense clouds. This decoupling is reinforced by
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A a (arcmin) Fig. 1 The inferred magnetic field directions for polarization measurements in the Galactic center are displayed on 850 jim contours from SCUBNJCMT (Pierce Price et al. 2000). Region I shows 60 pm polarimetery of the Sickle (Dotson et al. 2000). Region TI shows 100 pin polarimetry of the Arched Filaments (Dotson et al. 2000). Region 111 shows new 350 pm inferred magnetic field vectors along with the 350 pm vectors from (Novak et al. 2000). Important dust features are shaded and labeled. The axes scales are offsets in arcminutes from the position of Sgr A* (02000 = 17”45”’40504,6~000 = -29”00’28!’07).
noting that in Figure 1, the magnitudes of the 350 p m polarization vectors tend to be lower in regions of high submillimeter flux. Though the field measured by Hertz is not local to the individual clouds, the dense clouds in the central
30 pc do interact with the field. One example of this interaction is the region of the Sickle (See Fig. 1) Here, the magnetic field direction lies in a toroidal orientation parallel to the long axes of the molecular cloud. This direction is perpendicular to that of the poloidal field implied by the Radio Arc (see Fig. 2). Such an orientation is indicative of a field that has been sheared by the cloud’s motion.
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-5 15
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A a (arcminutes) Fig. 2 Inferred B vectors are superposed on a 20 cm map (Yusef-Zadeh,Moms, & Chance 1984) taken with the VLA. Important thermal and non-thermal structures are labeled. The 100 pm vectors appear to trace the Arched Filaments. Note that the field in the molecular cloud associated with G0.18-0.04 is perpendicular to the field associated with the Radio Arc.
Similar shearing is observed in M-0.13-0.08 (Fig. I). It has been noted by Novak et al. (2000) that the flaring of the field near the southern edge of this cloud indicates that the field in this region initially had a different configuration but was sheared by the inward motion of M-0.13-0.08 into a direction parallel to the cloud’s trajectory. Along these same lines, the effect of the motion M+0.07-0.08 on the magnetic field lines is noticed at the southwestern edge of the cloud. The magnetic field lines are observed to trace the cloud edge indicating that the cloud is moving towards the region of low flux and poloidal field to the northeast of M-0.02-0.07. In the process of moving through a low density region threaded by a poloidal field, this cloud is “sweeping up” poloidal magnetic flux and shearing it into a toroidal configuration.
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The toroidal magnetic field structures in dense molecular regions OF the Galactic center indicate that the gravitation energy density in these regions is strong enough to overcome the magnetic energy density such that the fields (which are frozen into the matter) are slretched out by the motion of the clouds. This relationship is quantified in Figure 3. Here. the absolute deviation from a poloidal field is plotted againsl the 350 pm SHARC flux, and we find that for high fluxes, the field is in general toroidal, while for low fluxes, the field is poloidal. 100
80
60
40
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C
-2c
50
100
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F (Jy/SHARC beam) Fig. 3 The absolute value of the deviation of each ineasurement from a poloidal field is plotted against 350 p m flux in a 15" SHARC beam. Here the angles are used from all of the GC polarization measurements in Figure I. The assumption is that the polarization angle will not change significantly from 60 pm to 350 pm.
This result can be interpreted in a way that reflects the model of Uchida et al. (1985) referenced in the introduction. We start with the assumption that the magnetic field in thc Galactic center was initially poloidal. In overdense regions, gravitational energy density can be strong enough to shear the field into a configuration that is toroidal (along the direction of the motions of the clouds). In underdense regions, the magnetic field energy density is high enough to maintain its poloidal structure againsL gravitational forces. This argument implies the existence of a critical density at which the magnetic energy density and the kinetic energy densily are equal. In this case, the following holds.
We assume this equilibrium occurs at a flux corresponding to a field angle of 45" with respect to the poloidal direction. Fitting a line to Figure 3 allows us to calculate the flux density associated with this
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Collapse along Field Lines
Molecular Material
I
I
I
Relativistic Electrons
Fig. 4 Molecular clouds can produce relativistic electrons necessary for the illumination of NTFs by the following process. In regions of low density, the molecular material is dominated by the magnetic field, and we observe a poloidal field (A). MHD allows for movement of material along the lines of flux. In this way, the material can form clouds and gravity can begin to compete with the magnetic field energy density (B). At this stage, velocities of the molecular material with respect to the poloidal field can distort the field. This process continues (C) as the poloidal fields become sheared into toroidal ones in the vicinity of the cloud. Finally, oppositely-oriented magnetic fields near the cloud centers will be forced into contact by gravity and will reconnect, thereby releasing energy that energizes relativistic electrons. These electrons spiral along the external field and produce synchrotron radiation that we observe as an NTF.
angle to be 125 Jy/SHARC beam. This flux can be converted into a mass density assuming a line-of-sight dimension and a typical cloud velocity. This leads to a characteristic magnetic field strength estimate.
4
Conclusions
This estimate of the magnetic field strength is in agreement with that obtained by Yusef-Zadeh and Morris (1987) for the field strength i n the Radio Arc. O u r data seem to suggest that m G fields penetrate much of
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the inner 30 pc of the Galaxy, but that conditions there are such that in regions of high density, the fields are sheared by the gravitational motion of clouds. The idea that the field measured in the cold dust may be part of the same global field that is seen in the NTFs suggests a possible source of the relativistic electrons required to “light up” the filaments. Magnetic reconnection has been posited as the source of the energy for the relativistic electrons (Serabyn & Gusten 1991); however, it is most often invoked as a cloud-flux tube collision in which the magnetic fields of a flux tube become bent and magnetically reconnect. Figure 4 depicts an alternative scenario derived from the idea that poloidal and toroidal fields are part of a global magnetic field. Here, we have a situation in which some underdense molecular material is threaded by a poloidal field. In this case, the field is strong enough to maintain its structure against the relatively weak influence of gravity. However, the material is free to collapse along the field lines. If this occurs, the gravitational energy density can begin to overcome the magnetic pressure and the field will begin to be sheared into a toroidal configuration (parallel to the long axis of the cloud.) As this shearing continues, oppositely-oriented toroidal magnetic fields are brought into contact with one another. At this point, magnetic reconnection can occur releasing energy that can produce the relativistic electrons necessary for NTF formation. Such a scenario fits observations of the Sickle (G0.18-0.04) region well. It can be seen that the toroidal magnetic field vectors are coincident with the molecular cloud shown in Figure I . Also, the poloidal filaments of the Radio Arc extend outward from the Sickle in both directions (see Fig. 2) showing a similar relationship between a set of NTFs and a molecular cloud as exhibited in Figure 4D.
References Chuss, D. T.. Novak, G., Hildebrand, R. H., Dowell, C. D., Vaillancourt, J. E., Davidson, J. A., & Dotson, J. L. 2003, ApJ, accepted Dotson, J. D., Davidson, J., Dowell, C. D., Schleuning, D. A., & Hildebrand, R. H. 2000, ApJS, 128, 335 Dowell, C. D., Hildehrand, R. H., Schleuning, D. A., Vaillancourt, J. E., Dotson, J. L., Novak, G., Renbarger, T., & Houde, M. 1998, ApJ, 504,588 Dowell, C. D., Lis, D. C., Serahyn, E., Gardner, M.. Kovacs, A,, & Yamashita, S. 1999, in ASP Conference Series 186: The Central Parsecs of the Galaxy, ed. H. Falcke, A. Cotera, W. Duschl, F. Melia, & M. Rieke, 453465 Killeen, N. E. B., Lo, K. Y., & Crutcher, R. 1992, ApJ, 385,585 LaRosa, T. N., Kassim, N. E., Lazio, T. J. W., & Hyman, S. D. 2000, ApJ, 119, 207 Novak, G., Chuss, D., Renbarger, T., Griffin, G. S., Newcomb, M. G., Peterson, J. B., Loewenstein, R. F., Pernic, D., & Dotson, J. L. 2002, in press Novak, G., Dotson, J. L., Dowell, C. D., Hildebrand, R. H., Renbarger, T., & Schleuning, D. A. 2000, ApJ, 529, 24 1 Pierce-Price, D., Richer, J. S., Greaves, J. S., Holland, W. S., Jenness, T., Lasenby, A. N., White, G. J., Matthews, H. E., Ward-Thompson,D., Dent, W. R. F., Zykla, R., Mezger, P., Hasegawa, T., Oka, T., Omont, A,, & Gilmore, G. 2000, ApJ, 545, L121 Plante, R. L.,Lo, K. Y., & Crutcher, R. M. 1995, AaJ, 445. L113 Schleuning, D. A. 1998, ApJ, 493,81 I Schleuning, D. A., Dowell, C. D., Hildebrand, R. H., & Platt, S. R. 1997, PASP, 109, 307 Serabyn, E., & Gusten, R. 1991, A & A, 242, 376 Uchidd, Y., Shibata, K., & Sofue, Y. 1985, Nature, 317, 699 Yusef-Zadeh, F., Moms, M., &Chance, D. 1984, Nature, 310, 557 Yusef-Zadeh, F. & Moms, M. 1987, AJ, 94, 1178
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Astron. Nachr./AN 324, No. S 1, 18 1 - 187 (2003) / DO1 10.1002/asna.200385066
The Galactic Center Nonthermal Fi1aments:Recent Observations and Theory T. N. LaRosa”, Michael E. Nord*.j, T. Joseph W. Lazio’, Steven, N. Shore4, and Namir E. Kassim3
’ Department of Biological & Physical Sciences, Kennesaw State Univ, 1000 Chastain Rd., Kennesaw, GA 30144 USA
* Department of Astronomy, Univ. of New Mexico
’ Code 7213, Naval Research Laboratory, Washington, DC 20375-5351
‘ Dipartimento di Fisica “Enrico Fermi”, Universita di Pisa, 56100, Pisa, Italy Key words Non-Thermal Filaments
Abstract. The large-scale topology and strength of the Galactic Center magnetic field have been inferred from radio imaging of the nonthermal filaments (NTFs). These objects, which seem to be unique to the Galactic center, are defined by extreme aspect ratios and a high degree of polarization. Recent high resolution, wide-field VLA imaging of the GC at 90 cm has revealed new candidate NTFs with a wide rangc of orientations relative to the Galactic plane. We present follow up 6 cm polarization observations of 6 of these candidates and confirm 4 as new NTFs. Together the new 90 and 6 cm results complicate the previous picture of largely perpendicular filaments that trace a globally ordered magnetic field. NTF observations in general do not rule out any particular models for the origin of the NTFs. Hence we explore the idea that the NTFs are local, individual structures: magnetic wakes generated through the interaction of molecular clouds with a Galactic Center wind. Numerical simulations of the evolution of a magnetized wake will be discussed and compared with NTF observations.
1 Introduction Beginning with the discovery of the bundled nonthermal filaments in the Galactic Center Radio Arc (GCRA) by Yusef-Zadeh, Morris and Chance (1984) it was recognized that the Galactic Center (GC) is the site of some unusual magnetic phenomena. The identification of isolated nonthermal filaments (NTFs), at various locations within the central few hundred pc (Yusef-Zadeh & Morris 1985; Liszt, 1985; Balley & Yusef-Zadeh 1989; Gray et al 1991; Lang et al 1999; LaRosa, Lazio & Kassim 2001; Reich 2003) shows that magnetic phenomena are widespread within the GC. The NTFs are defined by extreme aspect ratios (10-100) and high intrinsic polarization (e.g., Morris & Serabyn 1996). Although they are unique to the GC, their origin and relationship to other phenomena there continues to excite considerable speculation (e.g., Heyvarts, Pudritz and Norman 1988; Uchida, Shibata & Sofue 1985; Chevalier 1992; Benford 1988, Rosner & Bodo 1992; Serabyn & Morris 1994, Shorc & LaRosa 1999; Chandran 2001; Bickncll & Li 200 1a). In this paper we present new 6 cm polarization observations of candidate NTFs revealed by recent advances in low frequency wide-field, VLA imaging of the G C (Nard et al2003a,b; LaRosa et al2000). We then place the new results in context by providing a general overview of the observed NTF characteristics. We do not find that the observations can rule out any particular NTF model. Finally, we describe our * Corresponding author: e-mail:
[email protected], Phone: + I 770423 6038, Fax: t I 770423 6625
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group’s efforts exploring the possibility that the NTFs are dynamic structures created by the interactions between molecular clouds and a GC wind.
2 Observations Recently improved wide-field, high resolution, VLA imaging of the GC at 90 cm (Nord et al2003a,b) has revealed a number of new NTF candidates. Their Figure 1 covering a few 100 pc, was constructed from a combination of VLA A & B array data giving a high resolution of 12”x 7”. Although it is not sensitive to large-scale low surface brightness features, this image shows all previously known NTFs except for G359.44+0.39 which was significantly resolved. The previously known NTFs are mainly perpendicular to the Galactic plane. However, the orientations of the candidate NTFs are more diverse, with several nearly parallel to the plane. To verify their status, we have made follow up VLA observations of several of these sources in 6 cm polarization in Oct of 2002 in CnB configuration in dual polarization mode with significantly better resolution (- 4”x 3”) than at 90 cm (for details see LaRosa et a1 2003). We concentrated on several candidates in the Sgr C region. Figure 1 shows the 90 cm AB subimage of Sgr C.
G359.22-0.16 is located about 20 pc in projection south of the Sgr C H I1 region. This source was detected at 18 cm (Liszt & Spiker 1995) and appears 5 pc long. Its surface brightness is not uniform: the peak occurs at the southern terminus where it’s slightly wider. The 90 cm morphology is similar. -0.3 ( S P).The Using the peak flux at 18 and 90 cm gives an estimate for the spectral index N 6 cm polarized intensity image of (3359.22-0.16 is shown in figure 2; it is 40% polarized, consistent with previous measurements of other NTFs. It is 7 pc long and roughly 0.5 pc wide. Based on this morphology, its polarization, and its nonthermal spectral index we confidently classify this source as an NTF. This is an important result since this is only the second confirmed NTF that is parallel to the Galactic plane. Furthermore, since the end of this source is less than 10 pc in projection from the Sgr C filament, if both are at the same distance they cannot be tracing a simple globally ordered field. There is considerable evidence that the magnetic field in the neutral medium along the Galactic plane in the Sgr A and Sgr B regions is parallel to the plane (e.g., Novak et al 2003a,b). (3359.22-0.16 could be related to this toroidal field, but SCUBA images (Pierce-Price et a1 2000) along the Galactic plane indicate that the thermal emission from the Sgr C environment is much less than in the Sgr B or Sgr A regions so Sgr C may not be dominated by the neutral medium. N
N
N
G359.43-0.13 lies northwest of the Sgr C filament. At 90 cm it has a distinctive X-shape. At 6 cm the X resolves into two, or perhaps, three filaments with significant curvature (see Figure 3). Unfortunately, our low signal to noise ratio prevented a detection of this system in polarization. Thus we hesitate to classify this system as an NTF although morphologically similar to other NTFs. Deeper observations of this system are warranted. G359.40-0.07, seen 5 pc south in projection from the Sgr C filament is the brightest NTF candidate at 90 cm. This source was also detected at 18 cm (Lizst & Spiker 1995) and we estimate a 18/90 cm -0.3. The faint extension of this source in the 90 cm image (G359.40-0.03) was spectral index of a not detectable at 6 cm over our integration time due to the intensity of the background emission along the Galactic plane. N
We also observed the three linear features in the Sgr B region, G0.39-0.12, G0.37-0.07, and G0.39+0.05 (figure 4) Only one, 60.39-0.12, was detected in 6 cm polarization at the 10% level. Their morphologies and similar orientation to the filaments in the GCRA strongly suggests they are NTFs. If so, they are the first N7Fs to be found north of the GCRA and extend the volume over which the NTF phenomenon occurs to over 300 pc along the Galactic plane.
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Fig. 1 90 crn subirn of the Sgr C region. Candidate NTFs are labelled
3 Review of NTF Characteristics We now consider these observations together with previous work to formulate a compendium of observed characteristics. The population: The region over which the phenomenon occurs is 300 pc x 75 pc on the sky. There are now 14 confirmed isolated NTFs and another dozen or so filaments in the GCRA. Two of the isolated NTFs are parallel to the Galactic plane. There are another 14 candidate NTFs. These are considerably shorter and have lower surface brightness. Several of these are parallel to the plane and a few are quite close in projection to other NTFs. It seems unlikely these are tracing the same globally ordered magnetic field. Magnetic fields aligned longitudinally: Rotation measure studies show that the NTF magnetic fields are aligned along the long axis of the filaments (Tsuboi et a1 1986; Yusef-Zadeh, Parastaran & Wardle 1997; Lang, Morris & Echevarria 1999). Estimates for the strength of the magnetic field from equipartition yield 100 pG (e.g, Gray et a1 1995)The rigidity of the filaments against the ram pressure of the surrounding highly turbulent molecular clouds suggest magnetic field strengths of 1 mG (Yusef-Zadeh & Morris 1987a,b).
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Subfilamentation and braiding (e.g., Liszt & Spiker 1995; Yusef-Zadeh, Wardle & Parastaran 1997): High resolution images of isolated NTFs invariably exhibit subfilamentation, and flaring at the ends. Moreover, the subfilaments cross and appear to be braided around each other. If the NTFs were tracing a global field a reasonable assumption is that this field will have relaxed to a potential or force free state but braiding suggests that the field can be tangled and that magnetic reconnection is taking place. The stability of such structures within a global field remains an open question. Peak intensity located near geometric center (LaRosa et a1 2000): This property may be related to the braiding. Where two filaments overlap the surface brightness must increase.
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Fig. 2 6 cm total polarkation of G359.22-0.16. The rcsolution is 3.75” x 3’’
5 . Nonthermal Spectra with curvature above 5 GHz: The isolated NTFs have 20/90 cm spectral indicies -0.45 a 2 -0.6 (e.g., LaRosa et al. 2000). However, they are very weak at 6 and 2 cm, showing a clear spectral turnover at these higher frequencies (Lang, Morris & Echevarria 1999). This is not the case for the bundled filaments: the NTFs in the GCRA have been observed up to 43 GHz and their spectra are flat, suggesting ongoing particle acceleration or extreme youth (Reich, Sofue & Matsuo 2000).
6. Association with molecular clouds: Nearly a11 the well studied NTFs appear to be associated with, and could be interacting with, molecular clouds (e.g., Serabyn & Gusten 199 1 Staguhn et a1 1998). One explanation for the origin of the high energy electrons illuminating the NTFs is particle acceleration driven by magnetic reconnection occuring at the interface between a global magnetic fieid and some local cloud magnetic field. Examples of such interactions are inferred from observations of molecular clouds at several locations along the GCRA (Serabyn & Morris 1994; Tsuboi, Ukitd & Handa 1997). While these data are interesting and compelling, there is no obvious brightness variation in the NTFs at the purported interaction sites. 7. Lengthwise variations of spectral index: If electrons are injected at one end of a filament, then given the synchrotron lifetime (determined by magnetic field strength) and the lengths of the NTFs one might expect to see some spectral variation. Lang, Morris & Echevarria (1999) found no significant variation in the spectral index as a function of length in both the northern and southern threads, and LaRosa et a1 (2000) studied the Sgr C filament and found a constant spectral index with position. For a 1 mG magnetic field the synchrotron lifetime of electron emitting at 20 cm is N 3 x lo4 years. If the electrons can stream at the local Alfven velocity, they can travel roughly 65 pc in this time. Since the NTFs are shorter than 65 pc, LY variation may not be expected but in the longest NTF, the Snake, the spectral index does vary (Gray et al 1995). At the locations of the kinks the spectrum is flatter, suggesting that acceleration is taking place there (see Bicknel & Li 20014. For G359.43+0.2
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Fig. 3 6 cm total intensity image of G359.43f0.13. The resolution is 3.86" x 2.8"
Fig. 4 90 cni subim of the Sgr B region. Candidate NTFs are labelled
LaRosa, Lazio & Kassim (2001) found that the 20/90 cm spectral index decreased un~formlywith the distance away from the Galactic plane. The most natural explanation is that the electrons are radiating in a weakening magnetic field. If so, thc length scale for the variation is only a few pc suggesting this system may be locul field.
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These characteristics do not, however, point to any dejinitive model for the NTFs and fundamental questions remain open. Are the NTFs tracing a global or local field? Are these structures static or dynamic? Why are the NTFs observed only at the Galactic center? Are the same physical mechanisms that generated the isolated NTFs also responsible for the filaments in the GCRA? What acceleration mechanism generates the high energy electrons? Is the acceleration local or is it distributed along the length of the NTF? Although a number of ideas have been proposed (for recent reviews see Morris 1998; Bicknel & Li 2001b) we confine our discussion to the cometary model advanced by Shore & LaRosa (1999).
4 Theory: The Comet Model The NTF phenomenon must be related to special conditions that arise at the Galactic center and not in the general interstellar medium. The Galactic center has the highest concentration of young massive stars in the Galaxy and there is growing evidence for bursts of star formation there (e.g., Simpson et a1 1999). It is therefore reasonable that the combined stellar winds and supernova explosions will collectively generate a large-scale Galactic wind. Continuum X-ray observations have detected hot thermal gas and spectral line measurements indicate that it is moving at velocities of a few thousand km s-l (Koyama et a1 1996). Recent evidence for a nuclear starburst and associated outflows is summarized in Bland-Hawthorn & Cohen (2003). If this wind advects a weak magnetic field and encounters an obstacle, such as a molecular cloud, the resulting cloud-wind interaction will generate a magnetic wake (Shore & LaRosa 1999). The magnetic field diffusion time through a cloud is orders of magnitude longer than the fow timescale so the flow wraps the magnetic field around the cloud, forming a long thin wake. Detailed 3-D numerical simulations of a molecular cloud-wind interaction (Gergori et al 2000) confirm this picture (see their figure 9) and show the stretching and amplification of the magnetic field. In this cometary model the NTFs are observed in projection and can therefore have any orientation with respect to the Galactic plane. The advantage of this model is that since the NTFs represent local amplification of a weak field, the total magnetic energy density in the GC is greatly reduced compared to a strong pervasive field. However, the viability of this model depends on the stability of the wakes that are produced. Dahlburg et a1 (2002) have numerically simulated the evolution of a 2-dimensional wake for Galactic center conditions. The initial conditions were an exterior wind speed of 2000 km spl, an external density of I w p 3 ,an external field of mG, and an interior field of 1 mG. The code, which follows a piece of plasma in time, was initialized with small perturbations to insure the unstable modes were naturally excited using a linear code to determine the fastest growing wavelength (essentially a nonlinear magnetic KelvinHelmholtz instability) with which to normalize the length scales for a full nonlinear calculation. The results are shown in their figure 7 where the quantity pD2 (which is proportional to the synchrotron emissivity ) is plotted as a function of time (note that since the plasma is flowing longer times translate into increasing length along the filament). It is clear the field begins to shred, and the wake expands until it ultimately joins with the surrounding flow. However, the shredding begins after the the plasma travels approximately 50 times it own width so it is possible to obtain structures with aspect ratios similar to the NTFs. Furthermore, the peak brightness of the filament should occur where instability begins to grow, downstream from the interaction site. This may explain why the peak brightness of the NTFs occur near the midpoint rather than at one end. The instability may also be important for driving magnetic reconnection and creating the conditions for significant particle acceleration, which we are now studying.
5 Conclusions Based on morphology and polarization there are now 14confirmed NTFs, in addition to those in the GCRA. Two of the isolated NTFs are parallel to the Galactic plane. There are also several NTF candidates that show curvature and are also parallel to the Galactic plane. The candidate NTFs are considerably shorter and less intense than those previously identified. Several are quite close, in projection, to larger NTFs and
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would be difficult to explain if the NTFs are tracing a globally ordered field. At present observations d o not strongty restrict the theoretical possihilites. Continuing observational work combined with more delailed theoretical modeling will be required to elucidate the nature of these unusual objects.
Acknowledgements We thank Russ Dahlburg and Giorgio Einaudi for their willing and capable collaboration. Basic research in radio astronomy at the Naval Research Laboratory is supported by the Office of Naval Research.
References Bally, J. & Yusef-Zadeh, F. 1989, ApJ, 336, I73 Benford, G. 1998, ApJ, 333, 73 Bicknell, G., & Li, J. 2001a. ApJ, 548, L69 Bicknell, G. & Li, J. 2001 h, PASA, 18, issue 4. 43 I Bland-Hawthron. J. & Cohen, M. 2003, ApJ, 582, 246 Chandran, B. D. G. 2001, ApJ, 562,737 Chevalier, R. A. 1992, ApJ, 397, L39 Dahlburg, R., Einuadi, G., LaRosa, T.N. & Shore, S.N. 2002, ApJ, 568, 220 Gregori, G.. Miniati, F., Ryu, D. &Jones, T. J. 2000, ApJ, 543, 775 Gray, A. D., Cram, L. E., Eckers,R. D. & Goss, W. M. 1991, Nature, 353, 237 Gray, A. D., Nicholls, J., Ekers, R.D. & Cram, L. E. 1995, ApJ, 448, 164 Heyvaerts, J., Norman, C. & Pudritz, R. E. 1988, ApJ, 330, 718 Koyama, K., Maeda, Y., Sonobe, T. Takeshima, T. Tanaka, Y., & Yamauchi, S. 1996, PASJ, 48, 249 Lang, C. C., Moms, M. & Echevarria L. 1990, ApJ, 526, 727 Lang, C. C., Anantharamaiah, K. R., Kassim, N. E., Lazio, T. J. W., & Goss, W. M. 1999, ApJ, 521, L41 LaRosa, T. N., Kassim, N. E., Lazio, T. J. W. & Hyman, S. D. 2000, AJ, 119, 207 LaRosa, T.N., Nord, M. E., Lazio, T. J. W. & Kassim, N. E. 2003, in preparation LaRosa. T. N., Lazio, T. J. W. & Kassim, N. E. 2001. ApJ, 563, 163 Liszt, H. 1985, ApJ, 298,128 I Liszt, H. & Spiker, R. 1995, ApJS, 98,259 Moms, M. 1998, in The Central Regions of the Galaxy and Galaxies, IAU Symposium 184, ed. Y. Sofue (Dordrecht: Kluwer Academic Publishers), p. 33 I Morris, M. & Serabyn. E. 1996, ARA& A, 34,645 Nord, M. E., Lazio, T. J. W., Kassim, N. E., Hyman, S. D., LaRosa, T. N. & Duric, N. 2003a, AJ submitted Nord, M. E., Brogan, C . , Hyman, S. D., Lazio, T. J. W., Kassim, N. E., LaRosa, T. N., Anantharamaiah, K. & Duric, N. 2003b, these proceeedings Novak, G., et al. 2 0 0 3 , ApJ, 583, L83 Novak, G., et al. 2003b, these proceedings Pierce-Price, D., et al. 2000, ApJ, 545, L121 Reich, W. 2003, A& A in press Reich,W., Sofue, Y. & Matsuo, H. 2000, PASJ. 52, 355. Rosner, R. & Bodo, G. 1996, ApJ, 470, L49 Serabyn, E. & Moms, M. 1994, ApJ. 424, L91 Serabyn, E. & Gusten, R. 1991, A&A, 242, 376 Shore, S. N. & LaRosa, T. N. 1999, ApJ, 521.587 Simpson, J. P.,Whitteborn, F. C., Cohen, M. & Price, S. D. 1999 in ASP Conf. Ser. 186, The Central Parsecs of the Galaxy, eds. H. Falcke, A. Cotera, W.J. Duschl, F. Melia, & M.J. Rieke (San Francisco: ASP) p. 527 Staguhn, J., Stutski, J., Uchida, K.I., & Yusef-Zadeh, F. 1998, A& A, 336. 290 Tsuboi, M., Inoue. M., Handa, T., Tabara, H.. Kato, T., Sofue, Y. & Kaifu, N. 1986, AJ, 92, 818 Tsuboi, M., Ukita, N. & Handa, T. 1997. ApJ, 481, 263 Uchida, Y., Shibata, K. & Sofue, Y. 1985, Nature, 3 17, 699 Yusef-Zadeh, F., Moms, M., & Chance, D. 1984, Nature, 310, 557 YLisef-Zadeh, F. & Morris, M. 1985, AJ, 90,251 I Yusef-Zadeh, F. &Morris, M. 1987a, ApJ, 322, 721 Yusef-Zadeh, F. & Moms, M. 1987h, AJ, 94, 1 128 Yusef-Zadeh, F., Wardle. M. & Parastaran, P. 1997, ApJ, 475, L119
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Astron. Nachr./AN 324, No. S1, 189- 195 (2003) / DO1 10.1002/asna.200385068
Interaction between the Northeastern Boundary of Sgr A East and Giant Molecular Clouds: Excitation Mechanisms of the H2 Emission Sungho Soojong Pak2, Christopher J. Davis3, Robeson M. Herrnstein4, T. R. Geballe’, Paul T. P. Ho4, and J. Craig Wheeler6
’ Astronomy Program in SEES, Seoul National University, Shillim-Dong, Kwanak-Gu, Seoul 151-742, South Korea ’ Korea Astronomy Observatory, Whaam-Dong, Yusong-Gu, Daejeon 305-348, South Korea Joint Astronomy Centre, University Park, 660 North A’ohoku Place, Hilo, HI 96720, USA Harvard-SmithsonianCenter for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA Gemini Observatory, 670 N. A’ohoku Place, Hilo, HI 96720, USA Astronomy Department, University of Texas, Austin, TX 78712, USA
Key words Galaxy: centre - ISM: individual(Sgr A East), molecules -infrared: ISM: lines and bands Abstract. We have detected the v = 1 i 0 S(1) (A = 2.1218 pm) and v = 2 + 1 S(1) (A = 2.2477 pm) lines of HZin a region between the northeastern boundary of Sgr A East and the giant molecular cloud (GMC) M-0.02-0.07, in the Galactic centre. The broad line widths, the low peak velocities relative to the molecular clouds, and the moderate Hz v = 2 --t 1 S(1) to v = 1 + 0 S(1) line ratios can be best explained by a combination of C-type shocks and fluorescence. The detection of shocked H2 in this region is clear evidence that Sgr A East is driving material into both the GMC M-0.02-0.07 and the northern ridge found by McGary, Coil, & Ho (2001).
1 Introduction Sgr A East has frequently been interpreted as a supernova remnant due to its shell structure and non-thermal spectrum (Jones 1974; Goss et al. 1983 and references therein; and see the more recent references in Maeda et al. 2002). Some recent research, however, has suggested that the energetics, size, and elongated morphology (3’ x 4’ o r 7 pc x 9 pc at d = 8.5 kpc) of Sgr A East cannot have been produced by a typical supernova (Yusef-Zadeh & Morris 1987; Mezger et al. 1989). The origin of Sgr A East still seems to be an open question and the required energy to produce it is a key parameter in this issue (Mezger et al. 1989). In principle, the energy of the explosive event can be directly measured by studying regions where Sgr A East is colliding with ambient interstellar material. By tracing the dynamics of molecular gas, an interaction between the eastern part of Sgr A East and the giant molecular cloud (GMC) M-0.02-0.07 (also known as the ‘50 k m sP1cloud’) has been inferred (Genzel et al. 1990; Ho et al. 1991; Serabyn, Lacy, & Achtermann 1992; Mezger, Duschl, & Zylka 1996; Novak 1999; Coil & Ho 2000). Recent observations of NH3(3,3) emission in the region show that Sgr A East impacts material to the north and west as well (see Fig. I , also McGary, Coil, & Ho 2001). As direct evidence of this interaction, several 1720 MHz OH masers, which are a good diagnostic of the continuous, or C-type, shock excitation (Frail et al. 1996; Wardle, Yusef-Zadeh, & Geballe 1999), have been detected along the southern edge of Sgr A East and to [he north of the circumnuclear disk (CND) (Yusef-Zadeh et al. 1996). * Correspondin&author: e-mail:
1eeshOastro.snu.ac.kr.Phone: +8242 865 3248, Fax: +8242 861 S610 @ 2003 WILEY~VCHVcrkig GmhH & Co KO&.. Weinhelm
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i
c
,%I
~J2(i'JOI
Fig. 1 Central 10 x 15 pc region of the Galaxy. Contours representing the velocity-integrated map of NH3(3,3) emission are overlaid on a 6 crn continuum image of Sgr A complex from McGary et al. (2001). The black dot at the centre of image is Sgr A* and the mini spiral is Sgr A West. The CND is traced by the brighter part of continuum surrounding them and Sgr A East is seen by the outer part extended to the boundary of NH3 contours. The GMC M0.02-0.07 lies to the east of this region. The dashed box at the northeastern edge of Sgr A East encloses the 90" x 27" region observed in Hz. The solid line in the box is Slit 9 from which the Hz spectra shown here were extracted. Letters mark the positions of OH( 1720 MHz) masers with error ellipses scaled up by a factor of 15 (Yusef-Zadeh et al. 1999a).
Wardle et al. (1999) and Yusef-Zadeh et al. (1999b, 2001) detected H2 line emission in regions where OH-masers have been detected and they interpreted the emission as thermal. It is therefore likely that Sgr A East is indeed driving shocks into the adjacent GMCs to the south and into the CND. However, the fields observed by Wardle et al. (1999) and Yusef-Zadeh et al. (1999b, 2001) are restricted to the vicinity of the CND and cover only some of the regions where interaction of the Sgr A East shell with surrounding material is expected. Before one can hope to estimate the energy released in the event that created Sgr A East, it is necessary to observe additional interaction regions in diagnostic lines of Hz.
2 Observations We observed the Hz v = 1 --t 0 S(1) (A = 2.1218p.m) and the Hz v = 2 1S(1) (A = 2.2477pm) lines at the 3.8 m United Kingdom Infra-Red Telescope (UKIRT) in Hawaii on 2001 August 3 and 4 (UT), using ---f
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Fig. 2 €12 v = 1 + 0 S ( l ) and Ha v = 2 1 S ( l ) spectra from six positions along Slit 9. Indicated positions are relative to (Y = 17”45”45?3, 6 = -28d58“’58”; 52000. Left panels show the H P v = 1 3 0 S(1) spectra. The right three panels present both the HZ v = 1 0 S(1) and Hz v = 2 + 1 S(1) spectra from the positions where H:! v = 2 4 1 S(1) emission is detected; these are averaged over 3.4” on the sky to improve the S/N ratios. The dotted lines are Gaussian fits to the observed line profiles. The spectra are not corrected for instrumental broadening. -i
the Cooled Grating Spectrometer 4 (CGS4; Mountain et al. 1990) with a 31 I/mm echelle grating, 300 mm focal length camera optics and a two-pixel-wide slit. The spatial resolution along the slit was 0.90” for HZ v = 1 --* 0 S(1) with the grating angle of 64.691 degree and 0.84” for H2 v = 2 -+ 1 S(1) with 62.127 degree, respectively; the slit widths on the sky were 0.83” and 0.89”, respectively, for these two 90”. Thc instrumental resolutions, measured from Gaussian fits to sky configurations. The slit length is lines in our raw data, were 17 krri s-’ for Hz v = 1 + 0 S( 1) and 19 kin s-’ for HZ v = 2 + 1 S( l ) , respectively.
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Position along Slit 9 (arcsec) Fig. 3 Derived line parameters for the spectra along Slit 9: (a) line center velocity; (b) line width; and (c) integrated line intensity. Indicated positions are as in Fig. 2; positive towards NE. Thc filled circles and the open circles represent the +50 km sC1 component and the 0 kin s-’ Component of Ha v = 1 --* 0 S(l) spectra, respectively, and the filled and open squares (the +50 km s-’ and 0 k m s-’ component) of the NH:3(3,3), from McGary et al. (2001). In the panel (c), the solid and dashed lines dcnote the total intensity of HZ v = 1 + 0 S ( l ) and NH:3(3,3), respectively, the latter scaled by The range of our Hz observation is denoted by two vertical lines. The decrease in NHs flux at positions greater than +30” is a result of reduced sensitivity at the edgc of the mosaic (McGary et al. 2001).
Ten parallel slit positions were observed, sampling a 90” x 27” area on the northeastern boundary of Sgr A East. The slit was oriented 40’ east of north for each measurement; adjacent slit positions were separated by 3” perpendicular to the slit axis. The coordinates at the centre of the observed area are (Y = 17h45m45?9, 6 = -28d59”’05” (52000) (see Fig. 1). Only the ninth slit position, hereafter called ‘Slit 9’, was observed in both Hz v = 1 + 0 S(1) and H2 v = 2 + 1 S(1), the line ratio of which we can use to constrain models for the excitation of Hz. Here we present the results of Slit 9, rather than the whole data set, as a preliminary report. We aim to concentrate on the excitation mechanism ofthe detected H:! emission.
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Bright (1.6 - 21 x 1U-I8 W n-’ arcsw-2) HZ emission was detected from most of the observed region along the northeastern boundary of Sgr A East. From the Hz v = 1 + 0 S ( l ) spectra in Fig. 2 we measure line centers and line widths along the interaction region. Each spectrum is well fitted by one or two Gaussian components. Fig. 3 show the distributions of the derived Hz v = 1 i 0 S(1) line parameters along Slit 9. For direct comparison, we include in these figures data from the NH3(3,3) observations of McGary et al. (2001); the NH3 emission essentially traces the cool ( 5 100 K), dense ( l o 5 cnir3)),cloud material. From these data we note the following.
-
In Fig. 3a there are two velocity components, at ~ . S R 0 krri srl and +50 kin s r l . Both components are evident in Hz and NH:%.The +50 kui srl component of the NH3 emission traccs the GMC M-0.02-0.07, while the 0 kin s-l component corresponds to the “northern ridge” of McGary et al. (200 I ). The variation in the velocity of either component is less than 20 krii s r l . N
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The H2 line widths in Fig. 3b are generally much broader than the NH3 widths. The NH:r component at +20 kiii s-’ seen in Fig. 3a at negative offsets seems to trace hotter gas according to the NH3(2,2) to (1,l) line ratio map by McGary et al. (2001), which they suggest may be the result of an impact by Sgr A East. The distribution of the total intensity of Hz (the solid line in Fig. 3c) is similar to that of NH:3 (dashed line). It should be noted that NH3 is attenuated at positions greater than +30” by the edge of the primary beam in the VLA mosaic (McGary el al. 2001). There is a small discrepancy between NH3 and H2 between t 1 0 ” and 1-30” which may be explained either by exhaustion of the source of excitation (e.g. shock energy or UV photons) or by obscuration, at the inner, more dense, regions of the cloud. The brightest emission along Slit 9, between offsets -20” to -45” (the southwestern part), arises from the outer Hz clumps of Yusef-Zadeh et al. (2001).
4 H2 excitation The Hz v = 2 + 1 S(1) line was detected at thrcc locations along Slit 9, at positions NE 23.”7, SW lo.”& and SW 22.”2 relative to the centre of the slit ( o = 17”45”’45?3, 6 = -28d58’1158s; J2000). From these data we measured line ratios Hz v = 2 --* 1 S(1) / v = 1 + 0 S(1) of 0.40 0.12, 0.51 0.17, and 0.27 i 0.07, respectively (see Fig. 2). At other positions only the HZ v = 1 + 0 S ( l ) line was detected, with 3 a upper limits to the ratio of 0.5.0.6, and 0. I at offsets of NE 31.”8, NE 4.”8, and SW 44.”7 along Slit 9, respectively. Fluctrescent excitation in a low-density PDR (n(H2) < 5 x lo4 cin-”) should yield a ratio of‘ about 0.6. A lower ratio is expected in a more densc PDR environment (Black & van Dishoeck 1987), or in a shock. There are two basic types of shock: ’*jump” or J-type and “continuous” or C-type (see Draine & McKee 1993 for a review). A J-typc shock is formed in a highly ionized or weakly magnetized gas. Fluid parameters such as density and temperature undergo a discontinuous change (jump) at the shock front where the molecules may be dissociated. J-typc shocks (with vclocities greater than about 24 kin s - ‘ ) will completely dissociate the molecules (Kwan 1977); H2 emission occurs from a warm, recombination plateau in the post-shock rcgion. J-type shocks typically produce low line intensities cotnparcd to C-type shocks and HP v = 2 1 S(1) / 7; = 1 0 S(1) line ratios as large as 0.5 arc possible (Hollenbach & McKee 1989). At lower shock velocities, below the H2 dissociation speed limit, J-type shocks may yield much lower line ratios; < 0.3 (Smith 1995). In a C-type shock, where the magnetic field softens the shock front via ion-magnetosonic wave propagation such that the fluid parameters change continuously across the $hock front, the H2 dissociation speed limit is much higher (- 45 kin srl: depending on the density and magnetic field strength in the pre-shock gas). Smaller line ratios of about 0.2 are then predicted (Smith 1995; Kaufman & Neufeld 1996).
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From the observed ratios alone we are not able to unambiguously distinguish between excitation mechanisms. Our results can either be explained by fast J-type shocks or dense PDRs, or by a combination of fluorescence and either C-type shocks or slow J-type shocks, since the higher line ratios associated with fluorescence will be tempered by the low H2 v = 2 + 1 S(1) intensities associated with collisional excitation in shocks. To help distinguish between the HZ excitation mechanisms, we consider kinematic information and the spatial variation of the line ratio along Slit 9 . At most positions in Fig. 3b, the Hg line widths are high; typically 40 - 70 kni s-', but as high as 120 kin s-' in some positions. This suggests shock excitation and turbulent motions in the gas and tends to exclude the pure fluorescence models, where the H2 line emission generally arises from the stationary gas at the edges of neutral clouds illuminated by Far-UV photons from early-type stars. In shocked regions the line ratio is found to be constant over a wide range of Hg v = 1 + 0 S(1) intensities and spatial positions (Davis & Smith 1995; Richter, Graham, & Wright 1995), although this is not necessarily predicted from theory (Draine & McKee 1993). Conversely, in a PDR the ratio is sensitive to the incident F W flux and the molecular gas density; the H2 v = 1 + 0 S(1) intensity increases but the Hg v = 2 + 1 S(1) / v = 1 + 0 S(1) ratio decreases with increasing gas density or UV intensity (Usuda et al. 1996; Takami et al. 2000). Thus an unchanging H2 v = 2 i 1 S(1) / v = 1 + 0 S(1) ratio is found in shocks, while a varying ratio is expected in the pure fluorescent case. The measured line ratio and the H2 v = 1 + 0 S( 1) intensity in Fig. 3c, show evidence of an anti-correlation in our data, as expected in dense PDRs. Although the wide line profiles point to shock excitation, fluorescence appears to play a significant role at at least some locations. Considering the kinematics further, we note that J-type shocks produce narrow lines that peak at the velocity of the shock, while C-type shocks produce broader lines which peak at the velocity of the preshock gas and extend up to the shock velocity. Fig. 3a shows that there are two velocity components that are similar in Ha and NH3. The H2 emission traces hot (- 2000 K) gas and the NH3 cool (5 100 K) gas. Thus, if we assume that shocks are driven by Sgr A East into adjacent molecular clouds, whose velocities are given by thc NHS data (M-0.02-0.07 at +50 km s-l and the northern ridge at 0 k m s-'), then fast J-type shocks are inconsistent with our results, due to the low peak velocities of the HS lines relative to the molecular clouds. In summary, then, the wide line profiles and low peak velocities indicatc C-type shock excitation. However, the high values of the line ratio at some positions along Slit 9 and the spatial variation in that ratio, point to a fluorescent component to the excitation in some locations. A combination of C-type shocks and fluorescence (see e.g. Fernandes, Brand, & Burton 1997) is therefore the most reasonable explanation for the Hz excitation. Our conclusion on the C-type shocks is consistent with the detection of 1720 MHz OH masers to the north of the CND and to the south of Sgr A East (Yusef-Zadeh et al. 1996). Very recently Karlsson et al. (2003) detected the 1720 MHz OH masers at two positions near our target region. For the fluorescence, the source of the UV radiation could be either nearby early type stars or J-type shocks. However, as noted above, we see no evidence of J-type shocks in our data. Also, we cannot establish whether nearby stars are the source of the UV flux due to the lack of information on where or how many early type stars there are in the region. Acknowledgements We give special thanks to Young-Sam Yu and Tae-Hyun Kim for their help with the observation. Fig. 1 is reproduced from Fig. 10 of McGary et al. (2001) by permission of the AAS. The United Kingdom Infrared Telescope is operated by the Joint Astronomy Centre on behalf of the U.K. Particle Physics and Astronomy Council. This work was financially supported by the BK21 Project of the Korean Government. TRG's research is supported by the Gemini Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., on behalf of the international Gemini partnership of Argentina, Australia, Brazil, Canada, Chile, the United Kingdom and the United States of America.
References Black, J. H., & van Dishoeck, E. F. 1987, ApJ, 322, 412
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195
Coil, A. L., & Ho, P. T. P. 2000, ApJ, 533, 245 Davis, C. J., &Smith, M. D. 1995, ApJ, 443, L41 Draine, B. T., & McKee, C. F. 1993, ARA&A, 3 I , 373 Femandes, A. J . L., Brand, P. W. J. L., & Burton, M. G. 1997, MNRAS, 290, 216 Frail, D. A., Goss, W. M., Reynoso, E. M., Green. A. J., & Otrupcek, R. 1996, AJ, I I I , 1651 210, 565 Genzel, R., Stacey, G. J., Harris, A. I., Geis, N., Graf, U. U., Poglitsch, A,, & Sutzki, J. 1990, ApJ, 356, 160 Goss, W. M., Schwarz, U. J., Ekers, R. D., & van Gorkom, J.H. 1983, in Proc. IAU Symp. 101, Supernova Remnants and Their X-Ray Emission, eds J. Danziger & P. Gorenstein (Reidel: Dordrecht), p. 65 Ho, P. T. P., Ho, L. C., Szczepanski, J. C., Jackson, J . M., Armstrong, J. T.,& Barrett, A. H. 1991, Nat, 350, 309 Hollenbach, D., & McKee, C. F. 1989, ApJ, 342. 306 Jones, T. W. 1974, A&A, 30,37 Karlsson, R., Sjouwerman, L. O., Sandqvist, Aa., & Whiteoak, J. B. 2003, A&A, 403, 101 1 Kaufman, M. J., & Neufeld, D. A . 1996. ApJ. 456, 61 1 Kwan, J. 1977, ApJ, 216, 713 Maeda, Y., Baganoff, F. K., Feigelson, E. D., Morris, M., Bautz, M. W., Brandt, W. N., Burrows, D. N., Doty, J. P., Garmire, G. P., Pravdo, S. H., Ricker, G. R., & Townsley, L. K. 2002, ApJ, 570,671 McGary, R. S., Coil, A. L., & Ho, P. T. P. 2001, ApJ, S59, 326 Mezger, P. G., Duschl, W. J., & Zylka, R. 1996, ARA&A, 7,289 Mezger, P. G., Zylka, R., Salter, C. J., Wink, J. E., Chini, R., Kreysa, E., & Tuffs, R. 1989, A&A. 209, 337 Mountain, C. M., Robertson, D. J., Lee, T. J., & Wade, R. 1990, in Proc. SPIE Vol. 1235, Instrumentation in Astronomy VII, ed. D. L. Crawford (Bellingham: SPIE), p. 25 Novak, G. 1999, in ASP Conf. Ser. Vol. 186. The Central Parsecs of the Galaxy, eds H. Falcke et al. (San Francisco: Astron. SOC. Pdc.), p. 488 1998, Richter, M. J., Graham, J. R., & Wright, G. S. 1995, ApJ, 454, 277 Serabyn, E., Lacy. J. H.. & Achtermann, J. M. 1992, ApJ. 395, 166 Smith, M. D. 1995, A&A, 296,789 Takami, M., Usuda, T., Sugai, H., Kawabata, H., Suto, H., & Tanaka, M. 2000, ApJ, 529, 268 Usuda, T., Sugai, H., Kawabata, H.,Inoue, M. Y., Kataza, H., & Tanaka, M. 1996, ApJ, 464, 818 Wardle, M., Yusef-Zadeh, F., & Geballe, T. R. 1999, in ASP Conf. Ser. Vol. 186, The Central Parsecs of the Galaxy, eds H. Falcke et al. (San Francisco: Astron. Sac. Pac.), p. 432 Yusef-Zadeh, F., & Moms, M. 1987, ApJ, 320, 545 Yusef-Zadeh, F., Roberts, D. A., Goss, W. M., Frail, D. A,, & Green, A. J. 1996, ApJ, 466, L25 Yusef-Zadeh, F., Roberts, D. A., Goss, W. M., Frail, D. A., & Grcen, A. J. 199Ya, ApJ, 512, 230 Yusef-Zadeh, F., Stolovy, S. R., Burton, M., Wardle, M., & Ashley, M. C. B. 2001, ApJ, 560, 749 Yusef-Zadeh, F., Stolovy, S. R., Burton, M., Wardle, M., Melia, F., Lazio, T. J. W., Kassim, N. E., & Roberts, D. A. 19Y9b, in ASP Conf. Ser. Vol. 186, The Central Parsecs of the Galaxy, eds H. Falcke et al. (San Francisco: Astron. Soc. Pac.), p. 197
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Astron. NachrJAN 324. No. S1. 197 -203 (2003) / DO1 10.1002/asna.200385032
Sgr A East and its surroundings - a view with XMM-Newton Masaaki Sakano*’x3,Robert S. Warwick**’,and Anne Decourchelle****
’ Department of Physics and Astronomy, University of Leicester, Leicester LEI 7RH, UK ’ Japan Society for the Promotion of Science (JSPS)
* CEAIDSMIDAPNIA,Service d‘Astrophysique, C.E. Saclay, 91 191 Gif-sur-YvetteCedex, France
Key words The Galactic Centre, supernova remnants, Individual: Sgr A East, X-ray, Plasma emission
We present an X-ray study of the Sgr A East region based on recent XMM-Newton observations. The spectrum of Sgr A East can be represented by a two-component thin thermal plasma model with temperatures of 1 and 4 keV, both of which have reached ionization equilibrium state. The abundance of iron is found to be higher in the central region of the nebula, with Z = 3 4 solar, than in the outer area for which Z 0.5 solar. On the other hand, the abundances of other elements appear uniformly distributed with Z 1. We also detect a weak fluorescent Ka line from neutral iron in the outer region of source. We discuss the nature of Sgr A East on the hasis of these new X-ray results.
--
1 Introduction Sgr A East is a radio bright non-thermal structure surrounding the Galactic Centre. The oval shell-like structure in the radio continuum supports the view that Sgr A East is a supernova remnant (SNR), SNR GO.O+O.O (Jones 1974; Ekers et al. 1983). On the other hand, alternative interpretations have also been proposed, for example it could be the remnant of an explosion associated with the central massive black hole Sgr A*, which is located within Sgr A East. There is quite good evidence for a physical interaction between Sgr A West and East (e.g. Yusef-Zadeh et al. 2000), although many of the details, such as the origin of Sgr A East and its past evolution, remain open questions. Hard X-ray imaging observations of the Sgr A region were first made with ASCA (Koyama et al. 1996) and then with SAX (Sidoli & Mereghetti 1999); they traced an extended distribution of X-ray emission originating from a thin thermal plasma. More recently this region has been observed by CkandmlACIS, at a spatial resolution of 0.5 arcsec (Baganoff et al. 2003a,b; Maeda et al. 2002, 2003; Park et al. 2003). The bulk of the radiated energy of a SNR appears in the X-ray band, hence X-ray observations are essential for the study of this kind of source. In this paper, we report the XMM-Newfon results for Sgr A East and its surroundings as derived from both X-ray imaging and spectral data. XMM-Newton observations benefit from the large effective area, good imaging capability (- 5 arcsec) of the XMM mirrors, and good energy resolution of the EPIC CCD detectors. With this capability, the quality of the spectrum and statistics of the image are remarkable. The XMM-Newton observation of the Sgr A region was made on 2001 September 4 as a part of the XMMNewron Galactic Centre Survey. Further details of this observation and the data screening are reported in Sakano et al. (2003a), whereas a general description of the survey is given in Warwick (2002,2003) and Sakano et al. (2003~).In this paper, we assume a value of 8.0 kpc for the distance to the Galactic Centre (Reid 1993). masQstar.le.ac.uk, Phone: +44 116 252 35 10, Fax: +44 116252 33 1 1 * * e-mail: rswQstar.le.ac.uk.Phone: +441162523517,Fax: +44116252331I . * * * e-mail:
[email protected],Phone: +33 1690843 84, Fax: +33 169086577. * Corresponding author: e-mail:
@ 2001 WILEY-VCH Verlag GmhH & Co KGdA. Weinhem
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2 X-ray images
-29 07 00
Radio Shell
k M M J174540-290.
J174540-29045
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4550
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17:45:30
17:46:00
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45:40
17:45:30
17:46:00
4S:SO
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Fig. 1 (a) XMM/MOS1+2 image of Sgr A East in the 2-9 keV band. The dashed line indicates the position of the radio shell of S g r A East and the solid lines show the main source and background regions used in the spectral extraction. The two concentric circles represent suh-regions referenced in Fig. 4a and b. The small circular region enclosed by the thick dotted line, which encompasses Sgr A*, is excluded in the spectral accumulation. (b) The hardness ratio image (4.5-9 keV/24.5 keV) overlaid with the 2-9 keV contour. Black regions are harder in X-rays. The position of an extremely hard (black) source XMM J174540-2940.5 is also indicated (see Sakano et al. 2003d, e for detail). (c,d) The He-like iron and sulfur Ka-line images, respectively, where the underlying continuum is subtracted. The 2-9 keV contour is superimposed. All the coordinates are in 12000. The Galactic Plane (bn = 0) is also indicated. Fig. la shows the X-ray image from the MOS1+2 cameras. The brightest spot in this image corresponds to the location of Sgr A’. The X-ray emission around Sgr A* is found to be extended. Note the time interval including the Sgr A* flare reported by Goldwurm et al. (2003a, b) was not included in our analysis. Sgr A* was reported to be resolved in the X-ray band by Chandra (Baganoff et al. 2003a), whereas the recent very deep Chandra observation found that Sgr A* is in fact slightly extended in the quiescent state (Baganoff et al. 2003b). At any rate, Sgr A* carries only about 10% of the total X-ray flux in the central 10” region (Baganoff et al. 2003a) except during flaring episodes. Taking the spatial resolution of XMM-Newton of -5 arcsec (for MOS; Jansen et al. 2001) into account, our result is consistent with the Chandra result.
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There is a further X-ray emission to the east of Sgr A*, extended nearly parallel to the Galactic plane. I t is clearly distinct in both the lower and higher energy bands. This X-ray emitting region is mostly contained within the radio shell of Sgr A East and thc molecular dust ring (see Maeda et al. 2002,2003; Park et al. 2003), which suggests a strong physical connection between the observed X-ray and radio emission. The whole X-ray emitting region of Sgr A East extends across 200 arcsec, whereas the e-folding radius of its core is 28 arcsec in the 2-10 keV band. The hardness image (Fig. Ib) shows that the regions to the north and south of Sgr A East in Galactic coordinates are enhanced in the soft X-ray band or, alternatively, that there is a narrow (< 3’ wide) strip along the Galactic plane with a harder spectral form. These features are also seen in the Ckundru image (Morris et al. 2003; Park et al. 2003). The scale height of the molecular clouds in the Galactic Centre region is 5’-9’ (Tsuboi, Handa & Ukita 1999; Sakano 2000), which corresponds to the full extent of our X-ray hardness picture. In addition the soft enhancement is significantly brighter than its adjacent regions. The most plausible interpretation would seem to bc that the soft X-ray enhancement represents an X-ray outflow. However, we note that the soft emission does not seem to emanate directly from Sgr A*, but from a region located a couple of arcminutes to the east of Sgr A * . Further discussion is found in Warwick (2002,2003). Fig. Ic and Id show narrow-band images corresponding to the He-like iron (6.7-keV) and sulfur (2.4keV) K a lines, respectively, where the underlying continuum has been subtracted using adjacent bandpasses and assuming an averaged spectral shape of the whole field of view. The 6.7-keV line is clearly more concentrated in the core of Sgr A East than the continuum (Fig. la). This implies that the core of Sgr A East is more abundant in iron, or possibly that it is higher in temperature, or the combination of the two. In contrast, the 2.4-keV line peak is located on Sgr A*, as is the peak in the continuum. This nature is quantitatively evaluated with the spatially-resolved spectral analysis described in the next section (Section 3.2). In these XMM-Newtondata, we further found more extended structure in both the 6.7-keV or 6.4-keV line, where the latter represents the fluorescent Ktu line from neutral or low-ionized iron. Detailed results are given in Warwick (2003) and Predehl (2003) in these proceedings.
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3 X-ray spectrum 3.1
Model fitting of the whole Sgr A East spectrum
We extracted the source spectrum from a 100”-radius region, but excluded Sgr A* and its immediate surroundings (within a 24 arcsec radius region) and the response of a bright soft point source. The background spectrum is taken from an adjacent sky region at similar galactic latitude to the source region (see Fig. I and Sakano et al. 2003a for detail). Fig. 2 displays the resulting background-subtractedspectra taken with pn and MOS1 and 2. The spectra show several strong emission lines. We fitted the spectra with a phenomenological continuum and many narrow Gaussian lines, and found that the centre energies of most of these lines, except for the 6.4-keV line (see the previous section) correspond to K u lines from highly ionized ions. Then we estimated the ionization temperature for each atom from the ratio of the K-line fluxes of the helium-like and hydrogenlike atoms. Fig. 3 summarises the derived temperatures. The temperature is found to vary significantly from element to element; for example, the ion temperatures of sulfur and iron are 1.1 keV and 4.0 keV, respectively. This implies that the spectrum consists of multi-temperature components. In fact, when we tried to apply a single-temperature thermal model to the spectrum, it was clearly rejected. We next applied a two-temperature thin-thermal plasma model with a common ahsorption column, allowing the abundances of silicon, sulfur, argon, calcium, and iron in both the plasmas to be free. The abundance of nickel was linked to that of iron. In Ihe best-fitting result (Table I ) a counts excess below 2 keV is still evident which lead us to consider the possibility of patchy absorption. We assume that a
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S g r A East (all; Pcfabs(ZVmekal+6.4))
,
t 1
i i
5
2
10
channel energy (keV)
Fig. 2 The simultaneous fitting of the spectrum of Sgr A East within a radius of 100". A two-component thermal plasma model is employed with metal abundances of Si, S, Ar, Ca, and Fe allowed to vary. The low energy spectrum is modified by a partially covering absorber.
0
22
LSi '
~
"
'
~
"
"
S
~
'
'
'
Ar
'
~
"
'
Ca
Fe
Fig. 3 Ionization temperatures for each atom. These are estimated from the line ratio of K a lines from helium-like and hydrogen-likeions, assuming a thin thermal plasma in ionization equilibrium (Mewe 1985). The temperature of silicon may be subject to some systematic error due to uncertainty of the continuum model.
certain ratio 1- t (where E is a free parameter) of the emission is absorbed by a column of 7 x 1022Hcm-2 (fixed). The fitting result is found to be significantly improved with a probability by the F-test of99.996% (x2/dof=400.4/321).Table 1 and Fig. 2 summarise the best-fitting results. The bulk of the systematic trend in the residual is removed by this approach (Fig. 2).
The best-fitting heavy absorption column and its covering fraction (E) were found to be 1 4 . 3 ~ 1 0 ~ ~ H cm-' and 93%, respectively. The Galactic Centre region is known to be full of molecular clouds of
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Table 1 The spectral fitting results Params. FIX( i J A;l(i) NII(ii)
Unit
- Data Whole r < 28" 0.93 ( 0 . 8 3 4 . 9 7 ) 0.93 (fixed) 14.3 ( 1.3.3-16.4) 17.4 (15.9-19.1) 7.0 (fixed) 7.0 (fixed) 4.34 (1.884.84) 3.20 (2.67-3.79) 1.44 (1.16-1.73) 0.40(0.27-0.62) 0.97 (0.90-1.05) 0.86 (0.72-0.98) 13.0(9.8-17.0) 5.3(3.4-9.l) 2.5 (0.80-6.3) 1.7 (1.1-2.7) 1.7(1.1-2.6) 1.38(1.15-1.65) 1.14(0.83-1.47) 1.1 (0.52-1 7 ) 1.7(1.0-2.4) 2.12(1.68-2.57) 3.5(2.8-4.5) 1.25(1.12-1.40) 0.2(0.0-0.4) l.Z(O.8-1.8)
Whole -
[H cm-'1 [H cm-']
13.7 (13.5-14.5) ~
4.34 (4.194.48) 1.38 (1.34-1.41) kT,(2) [keV] 0.93 (0.86-1.00) Norm(2) 11.2(9.h-13.4) Z Si 4.8 (3.7-5.4) ZS 1.71 (1.541.88) z.4 r 1.26 (0.97-1.55) ZCi( 2.55 (2.14-2.97) &, .N2 1.34(1.27-1.41) Fe-Ke 1.5(1.0-1.9) [10-'ph s-' ~ r n - ~ ] X'ldof 422.11322 Fx (MOS) 12.2 [IO-"erS s-' ~ m - ~ ] k?;(l) Norm( 1 )
lkeV]
l . = y -8
-60"
r > 60" 0.91 (fixed) 13.2 (12.414.2)
0.4(0.0-0.5)
7.0 (fixcd) 4.80 (3.94-6.04) 0.64 (0.49-0.80) 0.90 (0.79-1.01) 5.8(4.1-8.1) 1.3 (0.7-2.2) 1.32(1.03-1.71) 1.05 (0.59-1.63) 2.4(1.7-3.2) 0.47(0.38-0.61) 0.7(0.4-1.1)
258.61238 3.9
190.34/198 4.8
1.4(1.1-1.8)
235.0/180
400.4132I 12.2
ti
0.93 (fixed) 16.0 (14.7-17.3) 7.0 (fixed) 5.5 (4.6-6.4) 0.31 (0.24-0.42) 1.01 (0.92-1.16) 5.6(4.3-7.8) 2.8 ( 1 . 6 5 . 8 ) 1.4(1.0-2.1) 1.1 (0.69-1.6) 1.9(1.4-2.6)
3.3
The unit for the normalisation is lo-'' / ' n,.+n1~dI~/(47r1l2), where n,, and Z L I I are the electron and proton number densities (crK3)and D is the distance of the source (crn). The fluxes are for the 1-10 keV energy band. The quoted uncertainties are at 90'% confidencc for one interesting parameter. various sizes from radio and far infrared observations. This patchy-absorption model with this large covering fraction (i.e. small fraction for the 'hole') must, therefore, be realistic enough. We also tried a three-temperature model, including a third additional lower temperature component as an alternative to the partial absorption model but this gave neither good improvements in the fitting nor a realistic soft X-ray luminosity. Thus we found that: ( I ) temperatures of the two components are 1 keV and -4 keV, which further confirms the estimate from the line-ratios (see Fig. 3); (2) the lower-temperature component has an orderof-magnitude higher emission measure than the higher-temperature one; (3) metal abundances are a few, to several tens, percent higher than the solar value on average, except for calcium, which is significantly more abundant (-2 solar) than the other elements. N
3.2
Model fitting of the spatially split spectra
-
a)
.t '"
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Fig. 4 The spectral fitting results for different sub-regions in Sgr A East. We have applied a two-temperature thin thermal plasma model plus a patchy absorption. (a) the core region with a radius of 28"; (b) thc annular region within radii of 28"-60"; (c) the whole source region outside of the 60"-radius circle (see Fig. la). For simplicity only the pn data are shown although the results apply to the simultaneous fitting of both the pn and MOS1/2 datasets.
The images of Sgr A East (Fig. Ic) show that the iron-line component concentrates in the core region more than continuum. Thus, we examined the spectral variation within Sgr A East. We extracted spectra
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from concentric regions of radii 28” (the core region) and 28/’-60” (the peripheral region) centred on the peak of the 6.7-keV-line image (Fig. lc). Also an outer region given by the full source region less the area within the 60’’ circle was also examined (see Fig. 1 a). The spectra from these three sub-regions were fitted with the final model defined in Section 3.1 with the covering factor for the absorption fixed, but all the other parameters free. Fig. 4 shows the resulting best-fitting models and Table 1 lists the result. We found that the iron abundance is higher for the smaller “core” region with Z = 3.5, whereas 2 M 0.5 in the outer region. The abundances of other metals do not significantly vary within Sgr A East from the solar norm. The temperatures of the plasmas are slightly lower in the core region. Neither the ratio of the intensity of the lowerand higher-temperature components varies. The 6.4-keV iron fluorescent line is significantly detected only in the outer region with an equivalent width of -75 eV. 3.3 Plasma parameters From the apparent extent of Sgr A East in the hard X-ray band in a 28” radius, we calculate the total plasma volume V to be 1 . 6 ~ 1 cm3, 0 ~ ~assuming a spherical shape. The plasma is found to comprise two components and we assume that each component exists separately in pressure balance with the other. Using the best-fitting parameters for the core region and introducing the total filling factor qttot,which is the sum of the filling factors of the two components, we estimated the filling factor of = 0.5qtot and 7114 = 0.5qtot, the density of n , , ~ = 23qtot -112 and n c , H = 6.1r,tii12, a total energy of E = 1.5 x 1049,vtot, 112 erg, and a total X-ray emitting mass of 19.17;: M a , where the subscriptions of L and H mean the lowerand higher-temperature components, respectively.
4 DISCUSSION 4.1 Is Sgr A East a SNR? Mezger et al. (1989) claimed that Sgr A East may be multiple supernovae based on their derived total energy of 6 x lo5’ erg. On the other hand, our derived total energy of the hot plasma in Sgr A East is 1.5 x 104g17toterg, which i s clearly much smaller than the Mezger et al. value and even smaller than the nominal energy for a single SNR (- 10” erg), thus one SNR can easily account for Sgr A East. Since the plasma has already reached ionization equilibrium, this estimate does apply to the full thermal energy in the observable X-ray band. On the other hand, the estimated (observed) mass of 1 .9qtot f Ma is sufficiently large to account for a supernova. With an age of 8000 yr for Sgr A East (Mezger et al. 1989),this mass may originate as either the ejecta or swept-up interstellar material. The localisation of the iron abundance enhancement suggests that the mass is predominantly that of the ejecta. This value of the mass, as well as the total energy, suggests that a single SNR is the likely scenario for the origin of Sgr A East. The most remarkable characteristic of Sgr A East is the unusually high temperature of 4 keV for a SNR of this age. One possible scenario to explain this high temperature is that the shock has interacted with an ambient plasma already preheated to a temperature of several keV, perhaps due to past activity in Sgr A*. N
4.2 Origin of the 6.4-keV line We detected significant iron fluorescent line emission at 6.4-keV only in the outer region of Sgr A East with an equivalent width of -75 eV (Section 3.2). This outer region is nearly coincident with the observed dust ringhhell (Mezger et al. 1986, 1989). The column of this dust shell is estimated to be 3 x (Mezger et al. 1989). If we assume a simple 10% cm7712 based on an average density of lo4 isotropic morphology for the cloud, this column can account for the observed equivalent width of the 6.4keV line. Hence the detection of the 6.4-keV line supports the idea that Sgr A East is actually surrounded by the dust shell. N
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The Metal Abundance in the Galactic Centre Region
The metal abundance in the Galactic Centre region remains uncertain with estimates ranging from near solar based on infrared star observations (Carr el al. 2000; Ramirez et al. 2000; Sellgren et al. 2003) to a few times the solar value or more based on a variety of other methods (e.g. Mezger et al. 1979; Murakami et al. 2001). Although there will inevitably be some contamination of our spectra by point sources which may have power-law like spectra and which are not spatially resolved with XMM-Newton, the impact must be small since the plasma emission dominates this region (Maeda et al. 2002). Our results should therefore be robust and show that, except for the core of Sgr A East, the iron abundance in the high-temperature interstellar matter that pervades the Galactic Centre region, is near to the solar value or possibly sub-solar (see also Sakano et al. 2003b). Acknowledgements M. S. acknowledges the financial support from JSPS.
References Baganoff, F. K., et al. 2003a, ApJ, submitted (astro-pW0102151) Baganoff, E K. et al. 2003b, these proceedings Carr, J. S., Sellgren, K., & Balachandran, S. C. 2000, ApJ, 530, 307 Dogiel, V. A., lchimura, A., Inoue, H., & Masai, K . 1998, PASJ, SO, 567 Ekers, R. D., van Gorkom, J. H., Schwarz, U. J., & Goss, W. M. 1983, A&A, 122, 143 Goldwurm, A,, et al. 2003a, ApJ, 584, 751 Goldwurm, A., et al. 2003b, these proceedings Jansen, F., et al. 2001, A&A, 365, L1 Jones, T. W. 1974, A&A, 30,37 Koyama, K., Maeda, Y., Sonobe, T., Takeshima, T., Tanaka, Y., & Yamauchi, S., 1996, PASJ, 48, 249 Maeda, Y., et al. 2002, ApJ, 570, 67 I Maeda, Y., et al. 2003, these proceedings Mezger, P. G., Penkonin, V., Schmid-Burgk, J., Thum, C., &Wink, J. 1979, A&A, 80, L3 Mezger, P. G., Chini, R., Kreysa, E., & Geniiind, H. -P. 1986, A&A, 160, 324 Mezger, P. G., Zylka, R., Salter, C. J., Wink, J. E., Chini, R., Kreysa, E., & Tuffs, R. 1989, A&A, 209, 337 Mezger, P. G., Duschl, W. J., & Zylka, R. 1996, A&AR, 7,289 Murakami, H., Koyama, K., & Maeda, Y. 2001h, ApJ, 558, 687 Park, S., et al. 2003, these proceedings Reid, M. J. 1993, ARA&A, 31,345 Sakano, M. 2000, Ph.D. thesis, Kyoto Univ Sakano, M., Warwick, R. S., Decourchelle, A,, & Predehl, P. 2003a, in preparation Sakano, M., Warwick, R. S., & Decourchelle, A. 2003b, AdSpR, submitted Sakano, M., Warwick, R. S., & Decourchelle, A. 2003c, Proc. “Japan-Germany Workshop on Galaxies and Clusters of Galaxies”, p.9 (astro-pW0212464) Sakano, M., Warwick, R. S., & Decourchelle, A. 2003d, MNRAS, 340,747 Sakano, M., Warwick, R. S., & Decourchelle, A. 2003e, these proceedings Sellgren, K., 2003, these proceedings Sidoli, L., & Meregherti. S. 1999, A&A, 349, L49 Tsuboi, M., Handa, T., & Ukita, N. 1999, ApJS, 120, 1 Wang, Q. D., Gotthelf, E. V., & Lang, C. C. 2002, Nature, 415, 148 Wang, Q. D. 2003, these proceedings Warwick, R. S. 2002, Proc. New Visions of the X-ray Universe in the XMM-Newton and Chandra era., in press (astro-phl0203333) Warwick, R. S. 2003, these proceedings Zylka, R., Mezger, P. G., & Wink, J. E. 1990, A&A, 234, 133
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Astron. Nachr./AN 324, No. SI, 205 -210 (2003) / DO1 10.1002/asna.200385107
Chandra ACIS Imaging Spectroscopy of Sgr A East Y. Maeda", K. Itoh', F. K. Baganoff2, M. W. Bautz2, W. N. Brandt3, D. N. Burrows3, J. P. Doty2, E. D. Feigelson", G. P. Garmire', M. Morris4, M. P. Muno', S. Park-?,S. H. Pravdos, G. R. Ricker2, and L. K. Townsley'
' Institute of Space and Astronautical Science, 3-1 - 1 Yoshinodai, Sagamihara, Kanagawa, 229-85 10, Japan ' Center for Space Research, MIT, Cambridge, MA. 02139, U.S.A. '
Astronomy and Astrophysics, Penn State University, 525 Davey Lab., University Park, PA. 16802, U.S.A. Physics and Astronomy, UCLA, Los Angeles, CA. 90095, U.S.A. Jet Propulsion Laboratory, MS 306-438.4800 Oak Grove Drive, Pasadena, CA 91 109, U.S.A.
Key words Galaxy: center - ISM: clouds - X-rays: individual (Sagittarius A East) -X-rays: ISM Abstract. We report on the X-ray emission from the shell-like, non-thermal radio source Sgr A East located in the inner few parsecs of the Galaxy based on observations made with the ACIS detector on board the Chandra X-ruy 0h.vewutory. The X-ray emission from Sgr A East is concentrated within the central Y 2 pc of the larger radio shell. The spectrum shows strong Kcu lines from highly ionized ions of S, Ar, Ca, 2 keV, absorption column and Fe. A simple isothermal plasma model gives electron temperature 1x H cm-', luminosity 8 x lo3* ergs ti-' in the 2-10 keV band, and gas mass 279 M o with a filling factor 17. The plasma appears to be rich in heavy elements, over-abundant by roughly a factor of four with respect to solar abundances. Accompanied with filamentary or blob-like structures, the plasma shows a spatial gradient of elemental abundance: the spatial distribution of iron is more compact than that of the lighter elements. These Chundrn results strongly support the long-standing hypothesis that Sgr A East is a supernova remnant (SNR). Since Sgr A East surrounds Sgr A* in projection, it is possible that the dust ridge compressed by the forward shock of Sgr A East hit Sgr A* in the past, and the passage of the ridge may have supplied material to accrete onto the black hole in the past, and may have removed material from the black hole vicinity, leading to its present quiescent state.
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1 Introduction The center of our Galaxy embodies a rich variety of phenomena which create diverse complex structures that are visible to us over a broad range of wavelengths. The radio emission from the central few parsecs of the Galaxy has several components, including a compact non-thermal source (Sgr A*) thought to he associated with the central massive black hole, a spiral-shaped group of thermal g a s streams (Sgr A West) that are possibly infalling to Sgr A*, and a 3 ' 5 x 2 ' 3 shell-like non-thermal structure (Sgr A East) (Ekers et al. 1975). S g r A East surrounds Sgr A* West in projection, but its center is offset by about 50" (2 pc). A number of arguments suggest that Sgr A* West is physically located very near or possibly embedded within Sgr A East. For the latter case, interaction between S g r A East and Sgr A* West would he inevitable, so S g r A East may h e a key for understanding the activity in the nucleus of our Galaxy (for a recent review, see Yusef-Zadeh, Melia & Wardle 2000).
*
Corresponding author: e-mail:
[email protected],Phone: 81 42759 8150, Fax: 81 42759 8455
@ XH!i WILEY-VCH V c i l q CimbH & Co KCPA, Weinhemi
(XKL4-61i71031SIOI-00016 17 SO+ 5010
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The Chandra X-ray Observatory (Chrmdru) with the Advanced CCD Imaging Spectrometer (ACIS) 10 keV and moderate spectral resolution with the detector combines the wide-band sensitivity up to superior high spatial resolution (0”.5-1”) of Chundru’s High-Resolution Mirror Assembly (HRMA). In this paper, we present a preliminary analysis of ACIS Galactic center observations with an exposure time of over 500 ksec, and review the first-epoch observations published by Maeda et a]. 2002. The other findings using the same data are separately reported in this volume by Baganoff, et al., Morris et al., Muno et al., and Park, et al.. XMM-Newton studies of Sgr A East are also presented by Predehl et al. and Sakano et al. in this volume. Throughout this paper we adopt a distance of 8.0 kpc to the Galactic Center (Reid 1993). N
2 Results 2.1 Images The first-epoch observation of the Galactic Center region was carried out early in the Chundru mission on 1999 September 21 over a period of 5 1 ks using the ACIS-I array of four abutted, frontside-illuminated CCDs. The follow-up observations reached 590 ksec as of 2002 June. During most of the follow-up observations, the central region of Sgr A East was pointed at the gap of the CCD array. To minimize any systematic uncertainties due to the gap, we used the first-epoch observations for an overall view of Sgr A East, while all the available data has been used to investigate the selected structures presented in this paper. The smoothed broad-band ( 1 5 - 7 keV) X-ray image of the Sgr A radio complex using the first-epoch observation is shown in Figure 1, where we have overlaid the X-ray image with radio contours from a 20 cm VLA image (F. Yusef-Zadeh. private communication). The outer oval-shaped contours are due to synchrotron emission from Sgr A East (Ekers 1983). The thermal radio source Sgr A West, an HI1 region located in the central parsecs of the Galaxy, appears on the western side of the Sgr A complex. Several bright compact X-ray sources can be seen in the images. With the follow-up observations, the number of the point sources reached over 2300 point sources within -20 pc (Muno et al. 2003). One of these sources is coincident to within 0”.35 with the radio position of the compact non-thermal radio source Sgr A* (Baganoff et al. 2001). In addition to the compact sources, the Sgr A West region shows bright diffuse X-ray emission superposed on a broader region of diffuse emission which peaks 1’ east of Sgr A*, and which appears to fill the central N 2 pc of the Sgr A East radio shell (Figure 1). Based on its spatial properties, we identify this diffuse X-ray emission centrally concentrated within the radio shell as an X-ray counterpart of Sgr A East. Figure 2 shows raw images in both the total and iron-K (6-7 keV) bands without exposure corrections. A broader feature in the center of the Sgr A East shell is especially conspicuous in the 6-7 keV band, where the flux is dominated by iron-K line emission. A curious linear feature 0”.5 long, which we refer to as ’the Sgr A plume’ in Baganoff et al. (2001), can be seen extending (in projection) from the center of Sgr A East to the northeast. In addition to the Plume, we found several filamentary structures within the Sgr A East shell (Figure 2). The filamentary structures in the ACIS field of view is comprehensively reported in our companion papers Morris et al. (2003) and Park et al. (2003). Notably, neither significant X-ray continuum or line emission is seen from the location of the radio shell.
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2.2 Spectra By extracting counts from a circular region of radius 40” within the shell, we made an over-all spectrum as shown in Figure 3. The spectrum exhibits a continuum plus emission lines which give critical information on the physical state of the plasma. We first fit the spectrum with a thermal bremsstrahlung model having four Gaussian emission lines of unspecified energy, all absorbed by an interstellar medium having cosmic abundances. The emission lines can be attributed to the Ka: transitions of the helium-like ions of sulfur, argon, calcium and iron. The line width for each line is consistent with being unresolved. The existence of
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the highly ionized ions confirms the presence of an optically thin thermal plasma. The equivalent widths of the lines are relatively small for the first three elements (EW 0.1-0.2 keV) hut very large for iron (EW 2 3.1 keV). The continuum temperature is around 3 keV and the line-of-sight column density is around 1 x loz3cm312. The large equivalent width of the iron line suggests that the Sgr A East plasma is enriched in heavy elements. Since the high ionization state of the iron K-line can be reproduced by a plasma in collisional ionization equilibrium (Masai 1994), we fit the spectrum to models of an isothermal plasma having variable elemental abundances (MEKA; Mewe, Gronenschild & van den Oord 1985), modified by interstellar absorption. The model with solar elemental abundances (Anders et al. 1989) was rejected with xz = 309 ( 1 85 d.0.f.) because it does not reproduce the large equivalent width of the iron line. The best-fit (x2= 217 with 184 d.0.f) was obtained using heavy element abundances which are 4 times solar. These results are given in Table 2 and are shown as a solid line in Figure 3.
=
Table 1 Best-fit parameters to the Sgr A East spectrum fitted with the MEKA model
Parameter [unit]
Best fit value
N H [loz2cmP2]
11.4(10.5-12.3)
kT, [keV]
2.1( 1.9-2.4)
Z
3.9(2.9-5.9)
Normalization
1.1 (0.9-1.3)
217( 184)
X2(d.o.f.)
/ ~ L , ~ L H/ (47rD2), ~V where 71, is the electron number density (cm-’)), ~ L H Normalization: is the proton number density (cm-’)). and D is the distance to the source (cm). n, = 1.8 x I L H for the best tit. 3 What is Sgr A East? The X-ray spectrum enriched by heavy elements suggests that the X-ray plasma is dominated by supernova ejecta. The small gas mass of 271; Ma and thermal energy lo4’ ergs are consistent with the ejecta by a single SN explosion. The X-ray images show filaments which are usually detected shocks in SNRs. Therefore, these results strongly supports the long-standing hypothesis that Sgr A East is a single SNR (Ekers et al. 1983). Rho and Petre (1 998) defined a new class of composite SNRs showing centrally concentrated thermal Xrays lying within a shell-like non-thermal radio structure. They called these objects “mixed morphology” (MM) supernova remnants and identified 19 members of the class. With the centrally concentrated X-ray cmission we find here, and its well-established non-thermal radio shell, Sgr A East becomes a ncw member of the class of MM SNRs. The non-thermal radio emission from Sgr A East in the direction of Sgr A West is heavily absorbed by Sgr A West (Yusef et al. 1987). This fact convincingly indicates that, along the line-of-sight, Sgr A West lies in front of the Sgr A East shell. However, the distance between Sgr A West and Sgr A East along the line-of-sight is uncertain. It is quite possible that they lie at nearly identical distances, in which case the front edge of the expanding Sgr A East shell has probably reached and passed through Sgr A*. This configuration is simply illustrated in Figure 5, see Baganoff t al. (2001) and Maeda et al. (2002) for a more thorough discussion.
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Fig. 1 Smoothed X-ray image (1.5-7.0 keV) taken during the first-epoch obscrvations, overlaid with 20 cm radio contours (white: Yusef-Zadeh private communication). Both exposure and vignetting corrections were applied. The field center is at Sgr A*, and the pancl size is 8’.4 x 8l.4.
Acknowledgements This work has been supported in parts by NASA contract NAS8-OI 128, NAS 8-38252 for the Chandra X-Ray Observatory.
References Anders, E., & Grevesse, N. 1989, Geochimica et Cosmochimica Acta, 53, 197 Baganoff, F. K. et al. 2001, ApJ submitted Baganoff, F. K. et al. 2003, in this volume Ekers, R. D., van Gorkom, J. H., Schwarz, U. J. & Goss, W. M. 1983, A&A, 122, 143 Maeda. Y. et al. 2002. ADJ. 570.671 Masai, K. 1994, ApJ, 437,770 Mewe. R., Gronenschild, E.H.B.M., & van den Oord, G.H.J. 1985, A&AS, 62, 197 Morris, M. P. et al. 2003, in this volume Muno, M. P. et al. 2003a, ApJ, Submitted Muno, M. P. et al. 2003b, in this volume Park, S. et al. 2003, in this volume Predehl, P. et al. 2003, in this volume Reid, M. J. 1993, Annual Review of Astronomy and Astrophysics, 3 I , 345 Rho, J . & Petre, R. 1998, ApJL, 503, L167 Sakano, M. et al. 2003, in this volume Yusef-Zadeh, F., Melia, F. & Wardle, M. 2000, Science, 287, 85
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Fig. 2 Deep Raw X-ray images in the total (left) and iron-K band (right) with 20 cm radio contours (green). Neither exposure nor vignetting corrections were applied. The dark lanes which crosses at the center are due to the gap of the CCD array.
L
1
2
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Energy (keV) Fig. 3 The ACIS-I spectrum of Sgr A East. Error bars are 1 0.The solid line corresponds to the best-fit value with the optically thin-thermal model summarized in Table 2. Fit residuals are shown in the bottom panel.
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data and folded model grp5Q/dinusaSmarga_grp.pha
I
1
data and folded model grp5Q/fillamant6_meq)b4rp.pha
I.
I.
I I
channel energy (kev)
1+_1IW
chonnsl energy (kev)
2 - c -
!W
Fig. 4 Spectra taken from two regions of Sgr A East. The left spectrum is taken from a central region of the shell, which shows the condensation at the iron-K band image in Figure 2. The equivalent width of the iron-K line was found as 5 keV. The right is for the filament A shown in Figure 2. The spectrum is dominated by thermal emission.
Fig. 5 A schematic diagram showing now a supernova remnant might regulate gas falling onto the supermassive black hole at the center of our Galaxy (Baganoff et al. 2001; Maeda et al. 2002). A supernova Sgr A East exploded about 10,000 yrs ago near the supermassive hlackhole in the Galactic Center. The inward shock wave heated up the ejecta that was detected in X-rays by ACIS, while the outward shock pushed gas towards the central black hole perhaps a thousand years ago.
Astron. Nachr./AN 324, No. SI, 211-215 (2003) /DO1 10.1002/asna.200385097
A Census of Dust Absorption at the Galactic Centre
',
Andy Adamson* Rachel Mason?, Emily Macdonald3, Gillian Wright2, Jean Chiar4, Yvonne Pendleton4, Tom Kerr', Janet Bowey', Doug Whittet6, and Mark Rawlings7
' Joint Astronomy Centre, Hilo, Hawaii 96720 U.S.A. ' Royal Observatory Edinburgh, Blackford Hill, Edinburgh EH9 3HJ, UK
'
University of Oxford, Department of Astrophysics, Keble Road, Oxford OX13RH, UK NASA Ames Research Center, Mail stop 245-3, Moffett Field, CA 94035, USA University College London, Dept. of Physics and Astronomy, Gower Street, London WClE 6BT. UK Rensselaer Polytechnic Institute, Department of Physics, Applied Physics and Astronomy, Troy, N Y 12180 Observatory, P.O. Box 14, University of Helsinlu, FIN-00014 Helsinki, Finland
Key words Interstellar: dust, molecules, extinction
Abstract. The Galactic Centre offers a uniquely valuable line of sight for studies of the nature of dust in the ISM, but the long-held assumption that the line of sight samples only the diffuse ISM has been subverted by ground-based and I S 0 observations demonstrating the presence of absorption bands due to solid-phase volatile ices in the field. Spatial variability of the observed features suggests a patchy distribution of even the foreground diffuse-medium absorption. To map this clumpy distribution and to produce an inventory of the dust components, we are carrying out a narrow-band imaging survey over a field which extends from the Galactic Centre to beyond the circumnuclear ring, using both IRCAM3 and Michelle on UKIRT to cover the 3 pm water ice, 3.3 pm PAH, 3.4 pm hydrocarbon and 9.7 pm silicate absorption features. This paper presents the rationale for this programme and reports on progress with analysis of the survey data.
1 Introduction
I
Through detailed infrared spectroscopy, we know much about the nature of the materials present in hydrocarbon dust grains in the diffuse ISM. The 3.4 pm absorption band, first seen in the diffuse ISM along the line of sight to the Galactic Centre (GC) (e.g. Willner et al. 1979, Butchart et al. 1986), has been thoroughly analyzed, and substructure within the band is identified with short-chained aliphatic hydrocarbons. I S 0 detections of the corresponding weak deformation modes at 6.85 and 7.25 pm (Chiar et al. 2000) have considerably narrowed the range of possible hydrocarbon materials (Pendleton & Allamandola 2002). The relationship of the diffuse-ISM hydrocarbons to other grain components is poorly known, although there is a correlation between the 3.4 p m band and the silicate feature over a range of extinctions (Pendleton et al. 1994, Sandford et al. 1995). Spectropolarimetry suggests that these materials cannot be physically associated with any form of aligned grains which are responsible for the visual and near-IR polarization of starlight (Adamson et al. 1999), and they cannot be associated with organic refractory mantles (i.e. a result of processing of ices - Chiar et al. 1998). The observations tend to favour models invoking small grains over those requiring carbonaceous mantles on large grains, though the key spectropolarimetry results will remain controversial until observations of organic and silicate polarization spectra can be obtained in exactly the same line of sight. In dark clouds, the 3.4 p m band is replaced by absorption at 3.47 pm (Allamandola et al. 1992), which is known to be associated with ices rather than the silicates (Brooke et al. 1996, Brooke et al. 1999, Chiar * Corresponding author: e-rnail:
[email protected], Phone: iOOl808969651I. Fax: +OO18089616516
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et al. 1996), indicating a major change in the nature of the hydrocarbon-bearing grains. Absorption due to polycyclic aromatic hydrocarbons (PAHs) has been suggested on the basis of a 3.25 pm absorption band seen in some dark cloud sources (Brooke et al. 1999, Sellgren et al. 1994). While seen in sources with deep ice absorption, correlations suggest that these materials are in fact more closely related to refractory grains such as silicates (Brooke et al. 1999). In fact, dense-cloud and diffuse-cloud PAHs may be different in nature: the dense-cloud feature at 3.25 pm peaks at a shorter wavelength, indicating a lower temperature and/or different composition, than the feature observed toward GCS3 in the Galactic Centre (which peaks at 3.28 pni - Chiar et al. 2000).
2 Rationale Our CGS4 spectroscopy of dust grains in the Galactic Centre (Chiar et al. 2002) permitted complete deconvolution of the complex of ice and dust absorption features along this line of sight. Fig. 1 shows the deconvolved hydrocarbon absorption feature in the 3.2-3.7 pm region, comprising the 3.4 /m aliphatic hydrocarbon band and a broad (FWHM 0.15 pm) absorption centered around 3.3 pm. Previous work, which has relied on fitting a local continuum across the sharp 3.4 pm band in what is a very difficult part of the telluric spectrum, has inadvertently removed this broad feature. The band is indicative of aromatic hydrocarbons, although it is too broad for absorption arising in PAHs embedded in solid grain material.
1
1 4.2
Fig. 1 Spectrum of hydrocarbon absorption in the Galactic Centre, deconvolved from the ice band (Chiar et al. 2002). Sharp negative lines are ratioing and atmospheric artifacts; m o w s indicate the main 3.4 pm band and the -3.3 pm “PAH’ absorption. Subfeatures near 3.5 p m are also associated with aliphatics.
The strength of the ice band varies dramatically across the Galactic Centre field. This may be due to additional ice absorption in the circumnuclear molecular ring (IRSS and IRS19 are both situated in or behind the ring, and these two sources have strong ice bands relative to the 3.4 pm feature; see the map of this region in Chiar et al. - this volume). Our spectroscopy shows that the 3.4 pm band also vanes across the field, raising the possibility that it arises in diffuse clouds which are detectably patchy, on very small transverse scales. The aim of the current programme is to determine the variability of hydrocarbon and ice absorption across the entire Galactic Centre, and so (i) to determine the scale size of the absorbing regions and (ii) further define the relationship between the aliphatic and aromatic hydrocarbon components.
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3 Observations We have carried out a 3 and 10 p m survey of the central 2 arcminutes of the Galactic Centre, in narrowband filters which sample all the key NIR dust features. The field of view includes the circumnuclear ring, and JHK imaging data already exist with which to constrain the extinction and short-wavelength wing of the ice absorption (Blum et al. 1996). Observations were obtained with the UKIRT facility imagers IRCAM3 and Michelle, in the summer of 2001 and 2002 respectively.
3.1 3 pm photometry IRCAM3 mosaics were obtained using the 3. I , 3.3, 3.4, 3.6 and 4.0 pm narrow-band filter sets. A mosaic in the broad L' filter was also obtianed. Fig. 2 shows a selection of images, and the filter bandpasses used. 100
:
3
32
34
36 Wsvelengthlw
38
42
a)
Fig. 2 a) The IRCAM3 filter set after convolution with atmospheric transmission; b) IRCAMS mosaics in the 3.1 & 3.3 p m filters. North is at the right, East is at the top
Fig. 3 a) L' continuum mosaic; b) Derived catalogue. Orientation is as in Fig. 2.
The L' mosaic image provides a deep L-band catalogue of the nuclear region, containing approximately 270 sources. Most of these are also detected in all the narrow-band filters. The basic catalogue is the deepest available wide-field L-band list of which we are aware in this region; only 10 of these sources have L-band photometry in the Blum et al. catalogue (Fig. 4), which covers a similar area of sky. By comparison,
A. Adamson et al.: Dust Absorption
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the existing near-infrared data required to provide additional model constraints is more extensive: 96 of these sources are detected in J H and K by Blum et al. Work is in progress on deeper JHK imaging data over the same field, taken with UFTI on UKIRT. This will permit constraints to be placed on extinction and thus remove an apparent degeneracy between the depth of the ice band and the degree of extinction.
4 ' 10
11
12
13
14
15
16
IRCAM L' instrumental magnitude
Fig. 4 Photometry within the overlap between our work and Blum et al. L band catalogue. Instrumental magnitudes from IRCAM3 are negative.
3.2
10 pm photometry
Michelle imaging in filter bandpasses in the 7.9,9.7, 1 1.6 and 12.5 pm filters was carried out in summer of 2002. A small chop was used (of order three arcseconds) to chop out as much as possible of the extended structure which pervades the Galactic Centre. The imaging is most sensitive in the 7.9 pm filter, where more than 20 sources are detected within the immediate environs of the Galactic nucleus. The majority of these sources are associated with stellar objects also detected in the L-band catalogue. These data will provide a measure of the silicate absorption complementing the hydrocarbon and ice measurements, for comparison with the data on organic and volatile components from the 3 pm survey.
4
Modelling
Modelling of the data consists of iterative fitting of a model spectrum of a hot stellar source, subjected to both continuum extinction and the combined effect of the silicate, ice and hydrocarbon absorptions, folded through the known bandpasses of the filters used. Since the images contain sources for which these features have already been observed spectroscopically, the scaling factors between observed counts and feature depths are self-calibrating. Initial results suggest that with the continuum extinction fixed by JHK photometry, reliable ice-band optical depths can be extracted, but to isolate the three main components of the hydrocarbon band from one another will be difficult. Clearly, silicate depths will be available only for a small subset of the sources detected in the L-band survey, but the detection of stellar sources in the silicate filters at this distance is a major challenge.
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Fig. 5 a) 7.9 pm image; b) 1 I .6 pm image. As for the IRCAM3 frames, North is to the right, East is to the top. Each detected source (the best detections are shown in boxes) shows the chophod signature, with positive beam in the centre and two negative beams on either side.
5 Summary Observations with the UKIRT facililty imager IRCAM3 have produced both an L-band catalogue, and narrow-band images over the central two arcminutes of the Galactic Centre. T h e catalogue is the deepest available over such a wide area in this region: some 270 sources are detected in the broad L' filter, an order of magnitude more than are present in the previous deepest list (Blum et al.). Narrow-band detections in filters distributed across the three-micron window are available for most of the detected sources. The fluxes are being fitted to produce a catalogue of the variation of key volatile and refractory dust components as a function of position within the Galactic Centre. Work is in progress on complementary JHK imaging, which supplements the J H K catalogue of Blum et al., and permits the continuum extinction to be constrained. Finally, w e have recently obtained Michelle observations over a similar field, detecting more than 20 faint point sources exterior to the central cluster, most of which are probably stellar in nature.
Acknowledgements We gratefully acknowledge the staff of UKIRT for their support of the telescope, and the Michelle team from the UK Astronomy Technology for their excellent instrument and its commissioning.
References Adamson,A.J., Whittet,D.C.B., Chrysostomou,A.C., Hough,J.H., Aitken,D.K.. Wright,G.S., Roche,P.F. 1999, ApJ, 5 12,224 Allamandola, L.J., Sandford,S.A., Tielens,A.G.C.M., Herbst,T.M. 1992, ApJ, 399, 134 Blum,R.D., Sellgren,K., Dep0y.D.L. 1996, ApJ. 470,864 Brooke,T.Y., Sellgren,K., Smith,R.G. 1996, ApJ, 459, 209 Brooke,T.Y., Sellgren,K., Geballe,T.R. 1999, ApJ, 517, 883 Butchart.1.. McFadzean,A.D., Whittet,D.C.B., Geballe,T.R. & Greenberg,J.M. 1986, A&A, 154, L5 Chiar,J.E., Adamson,A.J., Whittet,D.C.B. 1996, ApJ, 472, 66.5 Chiar,J.E.. Pendleton,Y.J., Geballe,T.R., Tielens,A.G.G.M. 1998, ApJ, 507, 281 Chiar,J.E., Tielens,A.G.G.M., Whittet,D.C.B., Schutte,W.A., Boogert,A.C.A., Lutz,D., van Dishoeck,E.F., Bernstein,M.P. 2002, ApJ, 537, 749 Chiar,J.E., Adamson,A.J., Pendleton,Y.J., Whittet,D.C.B., Caldwell,D.A., Gibb,E.L. 2002, ApJ. 570, 198 Pendleton,Y.J. & Allamandola ,L.J. 2002, ApJS, 138.75 Pendleton,Y.J., SandfordSA., Allamando1a.L.J.. Tielens,A.G.G.M., Sellgren,K. 1994, ApJ, 437, 683 Sandford,S.A., Pendleton,Y.J.,Allamandola,L.J. 1995, ApJ, 440, 697 Sellgren,K., Smith,R.G., Brooke,T.Y. 1994, ApJ, 433, 179 Willner, S.P., RusseIl,R.W., Puetter,R.C., Soifer,B.T. & Harvey,P.M. ApJ, 229, L65
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Astron. NachrJAN324, No. S I , 217-221 (2003) /DO1 10.1002/asna.20038S061
Thermal SiO observations of a shell attached to the nonthermal filaments in Sgr A Toshihiro Handa * I , Masaaki Sakano**2,and Masato Tsuboi***3
' Institute of Astronomy, University of Tokyo, Osawa 2-21-1, Mitaka, Tokyo 181-0015, Japan ' Department of Physics and Astronomy. University of Leicester, Leicester LE1 7RH, UK Astrophysics and Planetary Sciences, Ibaraki University, Mito, Ibaraki 310-8512, Japan
Key words Galactic Center, nonthermal filaments, molecular cloud, interstellar shock.
Abstract. The CS observations with NRO 4Sm telescope reveal that a dense molecular shell is located between Sgr A and the nonthermal filaments of the radio arc, and that the shell shows an interacting feature with the nonthermal filaments (Tsuboi et al. 1997). The shell i ociated with X-ray emission observed with ASCA and Chdndrd (Yusef-Zadah et al. 2002). These result suggests the shell is formed by a shock. The SiO lines in the v = 0 state are thermally excited and are thought to be a shock tracer. They are good probes to investigate the physical property of the shell. We observed the shell in SiO ( J = 1 0.71 = 0) and SiO ( J = 2 - 1, u = 0) using the NRO 4Sm telescope. Features associated with the shell clearly appear in both the SiO lines. The morphology of the shell in both SiO lines after adjusting the beamsize effect is similar in 2 - b - v 3-dimensional space. The intensity ratio of the two SiO lines ( J = 2 - 1 over J = 1 0) is uniform over the shell, suggesting the shell is uniform in density. Weestimate the average value of the ratio R ~ - l / l - uzz 0.9, which means the molecular &asdensity is about lo5 using multi-level population analysis with the LVG approximation. The morphology of the shell in SiO lines is quite similar to in the CS line in 1 - b - v 3-dimensional space. The intensity ratio of SiO ( J = 1 - 0) to CS ( J = 1 - 0) is almost uniform over the shell and Rs,o,cs Y 0.24. The ratio of the SiO (.I = 1 0) line over the CS ( J = 1 0) line is about 0.24, which means the relative abundance of SiO over CS, X(SiO)lX(CS) = 0.05 0.13 using multi-level population analysis with the LVG approximation. It ia close to the value found in an SNR shell. ~
~
~
~
~
1 Introduction The radio arc in the Galactic Center is a unique feature and its origin is still beyond the veil. Using spectral properties it is divided into two parts; the nonthermal filaments and the arched filaments. T h e nonthermal filaments are dominated by synchrotron emission and extend vertically to the galactic plane. The nontherma1 filaments are thought to b e a part o f the polari7ed lobe found by Tsuboi e t al. (1985). The arched filaments show thermal origin and connect Sgr A West and the nonthermal filaments. T h e morphology and polarized emission suggest the nonthermal filements were produced by a jet, like those seen in AGN, but n o point-like energy source was found on the nonthermal filaments. T h e three dimensional structure and origin of the radio arc are still behind a veil, although many investigations have been done. Millimeter astronomy shows that the Galactic center is rich in dense molecular gas. A CS line survey of the Galactic center region was made with the Nobeyama 45 meter telescope to explore the detailed structure of dense molecular gas (Tsuboi et al. 1999). O n e of the interesting features arising from the * e-mail: handaOi0a.s.u-tokyo.ae.jp,Phone: +81 422 345062, Fax: +81422 345041 * * e-mail: masOstar.le.ac.uk * * * e-mail: tsuboi O rnx.ibaraki.ac.jp
T. Handa et al.: Thermal SiO near the nonthermal filaments
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survey is an expanding shell contacting with the nontherinal filament. Morphological evidences show the shell is interacting with the nonthermal filament (Tsuboi et al. 1997). Recent high resolution images and spectra of the Galactic center region in X-ray require uniqueness of the shell. A large scale X-ray survey with ASCA shows an extended feature located at the shell and a composite spectrum in X-ray (Sakano et al. 2001). The X-ray feature is resolved as an association of X-ray sources and filaments with Chandra observations (Yuzef-Zadeh et al. 2002). The latest analysis shows that the X-ray spectra of the galactic eastern and western edges are different; the galactic western edge shows an ion line, but the galactic eastern edge shows a featureless spectrum. Even in the galactic western edge some clumps show highly ionized iron emission hut others show neutral iron emission (Bamba et al. 2002). These results suggest that the physical properties of the molecular gas in the shell, and the distribution of shocked molecular gas, are importanl. Thermal SiO emission is thought to be a shock indicator, because a shock produces much SiO gas from dust by evaporation and spattering.
~
~.~
~
-._l_
~
Fig. 1 The shell attached to the nonthermal filaments. Thick countours are radio continuum at 43 GHz (after Sofue et al. 1986) and a gray scale image with thin contourb is in CS at ULSR = 30 - 40 !un/s (Tsuboi et a1 1999) A CS feature marked by a thick arrow IS the shell
2 Observations The observations were made in April 2002 using the Nobeyama 45-meter telescope. Both SiO ( J = 1 - 0, v = 0) and SiO ( J = 2 - I , w = 0) lines are observed simultaneously. The beamsizes are 39.5” and 17.6”. and the mainbeam efficiencies are 0.8 1 and 0.50 at 43 GHz and 86 GHz, respectively. The observed area is 4’ 5 1 5 lo’, -10‘ 5 h 5 -4‘. The grid spacing is 20” or 0.82 pc assuming a distance to the Galactic center of 8.5 kpc. The antenna pointing was checked by five-point scans of nearby SiO maser sources, OH 2.6-0.4 and VX Sgr. Typical pointing difference between two pointing check procedures was 5”. We used acousto-optical spectrometers (AOS) with 124.7 kHz resolution. Although the original spectral resolutions correspond to 0.87 kmls and 0.44 kmls at SiO ( J = 1- 0,w = 0) and SiO ( J = 2 - 1,v = 0) frequencies, respectively, we smoothed and resampled to 5 kmls velocity resolution in each line.
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The line intensity was calibrated using the standard chopper wheel method. Typical rms noise levels in the atmospheric attenuation corrected antenna temperature scale are 0.1I K and 0.18 K in J = 1 - 0 and J = 2 - 1 lines, respectively. Liner or parabolic baselines were applied to all profiles between 'ULSR = - 150 krri/s and +150 km/s.
3 Results 3.1
Thermal SiO emission
Figure 2 shows the integrated intensity maps over ULSR = 15 40 km/s in SiO lines. The instensity scales are correced after main beam efficiencies. The observed half power beam width are shown in the left bottom corner of each panel. Besides the beamsize effects features in two lines are quite similar. The morphology in both lines is similar to that in CS ( J = 1 - 0) line (Figure 3). ~
,
1 (dcgrcc)
1 (degrcc)
Fig. 2 Integrated intensity distributions over ,uLsI< = 15 40 km/s in SiO J=l-O, v=0 (left panel) and SiO J=2-I, v=0 (right panel). The observed beamsize is shown at the left bottom corner in each panel. The intensity scale is corrected by the mdinbealn efficiencies. (The PDF version is in color.) ~
To avoid the beamsize difference in both transitions and pointing ambiguity, we convolve the image to be 69" beamsize by corresponding gaussians. Using the convolved images we made line intensity ratio maps of CS ( J = 2 1, u = 0 ) over CS ( J = 1 - 0) for every 5 k d s bin. The resultant channel maps show a rather uniform ratio and no significant features over 'L>LSR= 15 - 40 k d s . To check the uniformity of the line intensity ratio we made an intensity correlation diagram (Figure 4). The sample points are all pixels with 111,s~= 15 40 km/s by 5 k d s step over whole observed area by 20" grid. The majority of the pixels are along a straight line suggesting the ratio is uniform over the I - b - o space. The averaged ratio derived from the regression line is Xz-l/l-o = 0.88. Using a multilevel population calculation with the LVG approximation the ratio means typical density of the molecular gas is n(H2) = (1.0 - 1.7) x 10'ccrrir" at 40K 5 Tk 5 80K. The result is not affected by a beam filling factor if T ~ 2 0.1 ~ and~the factors ~ ~are the ~ same ~ in both , lines. The uniformity of the ratio suggests the gas density of the shell is rather uniform, although several pixels deviate from the primary component along the regression line. ~
~
T. Handa et al.: Thermal SiO near the nonthermal filaments
220
I (degree)
I (degree)
Fig. 3 Integrated intensity distributions over VLSR = 15 - 40 km/s in SiO J=I-O, v=O (left panel) and CS J=1-0 (right panel). The effective beamsize is shown at the left bottom comer in each panel.
3.2
Comparison with CS
The critical density of SiO ( 5 = 1 - 0 , v = 0) is similar to CS ( J = 1 - O), because the Einstein A coefficients of SiO ( J = 1 - 0 , =~ 0) and CS ( J = 1 - 0) are 3.0 x lop6, 1.8 x lop6, respectively. Therefore, the intensity ratio of these two lines shows the difference in abundance and indicates the shock strength. The SiO shell is seen over WLSR = 15 -40 km/s, and the shell in CS is seen over U L S R = 15 -45 km/s (Tsuboi et al. 1997). The overall features in SiO ( J = 1 - 0, II = 0) are quite similar to those seen in CS ( J = 1 - O), although the high velocity end of the integration range is different by single pixel in velocity (Figure 3). The SiO ( J = 1 - 0, v = 0) map integrated over U L S R = 15 45 km/s shows an additional component near 1 = 0.0105", b = -0.085'. This difference between CS and SiO may not be real; the integration range difference is only a single pixel in velocity and a strong emission ridge is located at the galactic northern edge of the shell at U L S R = 50 k d s both in CS and SiO. Another possiblity is that the difference is due to sampling grid between CS and SiO. Beside of this marginal difference the shell is also very similar in the position-velocity diagrams in both CS and SiO. After adjustment of beamsize to be 69" by gaussian convolution we made line intensity ratio maps of SiO ( J = 1 - 0.2) = 0) over CS ( J = 1 - 0) for every 10 k d s bin. The resultant channel maps show a rather uniform ratio and no significant features over VLSR = 15 - 35 k d s . To check the uniformity of the line intensity ratio we made an intensity correlation diagram (Figure 5). The sample points are all pixels in OLSR = 15 - 35 k d s by 10 k d s step over whole observed area by 20" grid. The majority of the pixels fall along a straight line, suggesting the ratio is uniform over the 1 - b - v space. The averaged ratio derived from the regression line is Rs,o/cs = 0.24. Using a multilevel population calculation with the LVG approximation the ratio is mainly affected by relative abundance between SiO and CS. The derived abundance is X(SiO)/X(CS) = 0.05 - 0.13 if X(SiO)/(dv/dr) = - lo-'". Comparison with typical values at L1157, a protostellar bipolar jet, ( w 3; Zhang et al. 2000) and at IC443G, an SNR shell, is ( w 0.06; Turner et al. 1992) suggests the shock strength of the shell is similar to that of a SNR shell. ~
22 I
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SiO( 1-0) Tinh [K]
Fig. 4 Intensity correlation diagram over U L S R = 15 - 40 km/s between SiO J=l-O, v=O and SiO J=2I , v=O. The mainbeam efficiencies are correcled in both lines. The beamsizes are adjusted to be 69’ after gaussian convolution.
CS( 1 -0)’I’mh
v( K
Lnds)
Fig. 5 Intensity correlation diagram over ULSR = 15 - 35 km/s between SiO J=1-0, v=O and CS J=l0. The mainbeam efficiencies are corrected in both lines and integration widths are the same as 10 k d s . The beamsizes are adjusted to be 69” after gaussian convolution.
Acknowledgements We thank Mr. Seiichiro Naito for his powerful assistance at the observations
References Bamba A. et al. 2002, in: New Visions of the X-ray Universe in the XMM-Newton and Chandra Era,, edited by E Jansen et al., ESA SP-488 (astro-ph/0202010) Sakano M. et al. 2002, ApJS, 138, 19 Sofue Y., Inoue M., Handa T., Tabara H., Hirabayashi H., Monmoto M., Akabane K. 1986, PASJ, 38,475 Tsuboi M., Inoue M., Handa T., Tabara H., Kato T. 1985, PASJ, 37, 359 Tsuboi M.. Ukita N., Handa T. 1997, ApJ 481,263 Tsuboi M., Handa T., Ukita N. 1999, ApJS 120, 1 Turner B. E., Chang K., Green S., Lubowich D. A. 1992, ApJ 339, I14 Yusef-Zadah F., Law C., Wardle M. 2002, ApJ, 568, L121 Zhang Q., Ho P. T. P., Wright M. C. H. 2000. AJ 119, 1345
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Astron. Nachr./AN 324, No. SI, 223 -227 (2003) / DO1 10.1002/asna.200385033
Absorption and Emission in the Four Ground-State OH Lines Observed at 18 cm with the VLA Towards the Galactic Centre R. Karlsson*
I,
Aa. Sandqvist', L.O. Sjouwerman2.', and J.B. Whiteoak4
' Stockholm Observatory, SCFAB-Albanova,SE-106 9 1 Stockholm, Sweden ' Joint Institute for VLBI i n Europe, Postbua 2, NL-7990 AA Dwingeloo, the Netherlands NRAO Array Operations Center, P.O. Box 0, Socorro NM 87801, USA ' CSIRO. Australian Telescope National Facility, Box 76, Epping NSW 2121, Australia Key words Galactic Centre, Sy A, OH Streamer, OH-masers, molecular clouds.
The OH distribution in the Sgr A Complex has been observed in the 1612, 1665, 1667, and 1720 MHz OH transitions with the Very Large Array (VLA), in the BnA and DnC configurations. Spectral line maps have been produced with a channel velocity resolution of about 9 km sC1, and with angular resolutions of 4" x 3". and 24" x 22", respectively. Some clear results are highlighted here: i) the existence of an OH streamer inside the Circumnuclear Disk (CND) near Sgr A*, iij absorption from the CND, iiij strong absorption towards the eastern and most of the western parts of the Sgr A East shell, iv) lack of absorption towards both Sgr A West and the compact H rr-regions to the east of Sgr A East, v) a double-lobed structure of the High Negative Velocity Gas (HNVG) oriented to the northeasl and southwest of Sgr A*, and vij compact point-like maser emission in all four transitions. In particular, a 1720 MHz maser at - 1 32 km s-' in the CND as counterpart to a 1720 MHz maser at +I32 km s-' in the CND, was observed.
1 Introduction The Hydroxyl radical, OH, is relatively abundant in the molecular clouds of the Galactic Centre (GC). The four hyperfine rotational transitions at 1612, 1665, 1667 and 1720 MHz are readily observable in absorption, although the 1612 MHz is often contaminated by interference from satellite navigation systems. In Local Thermodynamic Equilibrium, LTE, for an optically thin gas the relative intensities of the four hyperfine rotational transitions are 1 :5:9: 1. Studies of OH absorption therefore provide a useful tool for investigating the kinematics and physical environment in the GC. Maser emission is also observed at all four frequencies. The OH masers at 1612, 1665, and 1667 MHz are pumped by infrared radiation and associated with evolved stars and H 11-regions. The 1720 MHz OH masers are collisionally pumped by shocks penetrating into molecular clouds. Preliminary results, mainly concerning the O H Streamer at 1667 MHz, were reported in 1987 and 1989 (Sandqvist et al. 1987, 1989),and the full results of the BnA observations can be found in Karlsson et al. (2003).
2 Observations and data reduction The main lines at 1665 and 1667 MHz were observed with the Very Large Array (VLA) in its hybrid extended (BnA) and hybrid compact (DnC) configurations. The 1612 and 1720 MHz satellite lines were observed with the BnA configuration only. The observations were made in June 1986, and in October 1989. A total bandwidth of 3.125 MHz resulted in a total velocity coverage of about 550 km s-'. A brief * Corresponding author: e-mnil:
[email protected],Phone: 4 6 -
8 5537 8538, Fax: 4 6 -8 5537 8510
@ 2003 WILEY-VCH Verlag GnihH & Co. KGaA. Wcmhem
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Table 1 VLA 18-cm OH observational summary: BnA configuration, 18 antennas, R-polarization, June 1986
Frequency (MHz) 1612 1665 1667 1720
HPBW (arcsecxarcsec) 4.0 x 2.7 3.9 x 2.9 4.0 x 2.8 3.7 x 2.7
PA (")
55.3 64.7 61.1 57.1
Resolution (km sP1) 9.1 8.8 8.8 8.5
tintegration
(minutes) I44 169 173 148
Table 2 VLA 18-cm observational summary: DnC configuration,27 antennas, Dual L+R-polarization,October 1989
Frequency (MHz) 1665 1667
HPBW (arcsecxarcsec) 24.0 x 22.3 30.2 x 23.7
PA (")
Resolution (km s-')
29.8 46.3
8.9 8.8
tintegration
(minutes) 144 136
summary of the observations is shown in Tables 1 and 2. NRAO's Astronomical Image Processing System, AIPS, and VLA standard calibration procedures were used for the data reduction. The 1665 and 1667 MHz lines form one continuous data set because of overlapping velocity ranges. For the 1665 and 1667 MHz main lines, the continuum map was constructed by averaging some line-free channels on each side of the lines, before subtraction from the map cube. The 1612 and 1720 MHz continuum maps were produced by subtraction of a linear fit over frequencies from the visibilities in the u , %I plane. The maximum entropy method of AIPS was used for improving the quality of the maps shown in this paper.
3 Results The main results of this study are: i) full high spatial resolution line-of-sight velocity OH absorption maps in all four ground-state OH rotational transitions (Karlsson et al. 2003), ii) the OH Streamer (Fig. 1 and Fig. 2), iii) the double-lobed structure of the HNVG (Fig. 3), and iv) 10 new point-like OH maser sources (Table 3). The OH Streamer seems, in projection, to stretch inside the CND between the southwestern part of the CND and Sgr A*. Its length is about 2 pc and its density and velocity dispersion are highest in the "head', i.e. close to Sgr A*. The streamer is clearly detected between +76.5 km s-l and +23.6 !an sC1 in the 1612, 1665, and 1667 MHz lines. The head can be traced to a velocity of -20.4 km s-'. The OH Streamer is not detected at 1720 MHz. As the line-of-sight velocity drops, the streamer seems to shorten (Fig. 1). In the vicinity of Sgr A*, the head seems to turn in an anti-clockwise direction as the velocity decreases from +41.2 km s-l to +23.6 km s-'. At even smaller velocities, the streamer finally seems to coincide with Sgr A*. The velocity of the tail of the OH Streamer differs by more than 100 km s-' from the velocity of the CND at the position where the streamer appears to merge with the CND. This may imply that the OH Streamer and the CND are not located in the same plane. Figure 2 shows -TIIT, for the 1667 MHz line at +58.8 km sP1, where the streamer is seen at its full length. The general structure of the OH Streamer is retained in the -Tl/Tc map, although a quantitative analysis has to include a more detailed model of the background continuum (Karlsson et al., in preparation). Strong absorption is also observed towards both the eastern part, and parts of the western areas of the Sgr A East shell (Fig. 1 and Karlsson et al. 2003). There is a lack of absorption towards the compact H IIregions to the east of Sgr A East, and towards Sgr A West. Extended OH absorption is, however, detected
Astron. Nachr./AN 324, No. SI (2003)
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Fig. 1 1667 MHz OH absorption towards the Sgr A Complex. The contour levels are IS, 93 and 500 mJy/beam, and outline Sgr A East, the H ~iregionsfurther to the east, Sgr A West and Sgr A*, respectively. The velocities are indicated in the upper right corner. The OH absorption is shown i n grey-scale, and in colour in the electronic version of this paper. The OH Streamer is seen at full length between +76.5 km s-’ and +SO km s-’, and is very clear in colour.
in the region of Sgr A West with a maximum velocity around - 170 kni spl. Two lobes are observed, both in the line and -TIT, maps, to the northcast and southwest of Sgr A* (Fig. 3). This structure is attributed to the HNVG.
In the 1665 and 1667 MHz DnC observations, extended large-scale OH absorption is also observed towards the Radio Arc, the Bridge and the Sgr A Complex. Examples of absorption towards the “Pistol” and the “Bridge” are seen in Figures 4 and 5, at the velocities of -73.2 km spl and -99.6 km s-’ . The HNVG, the Expanding Molecular Ring (EMR), and the CND can all be identified in the full set of maps.
R. Karlsson et al.: 18-cm VLA observations of OH towards the GC
226
0
200
-28 58
-28
0
34
174236
Fig. 2 The 1667 MHz OH line-to-continuum distribution (-Tl/Tc),at +58.8 km SKThe I. OH-streamer is highlighteded by contours.
00
32 30 20 26 RIGHT ASCENSION(61950)
24
Fig. 3 The 1667 MHz OH absorption at - 170 km s-', showing the double-lobed, northwest to southeast, structure of the HNVG.The contours mark the same continuum components as in Fig. 1.
10
05
-28 40
45
-
50
Y)
:
I
55
c -I
vy
-2900
05
I1 30
00
41 30
Fig. 4 OH large-scale absorption at 1667 MHz in the Sgr A Complex at -73.2 km sK1, Observations made with the DnC configuration.
174400
4
Fig. 5 OH large-scale absorption at 1667 MHz in the Sgr A Complex at -99.6 km sK1. Observations made with the DnC configuration.
Emission is detected from OH masers in all four transitions. Table 3 lists the new detections found in the 1986 BnA data set, i.e. masers not previously published using other independent observations (1612 MHz: Lindqvist et al. 1992 and Sjouwerman et al. 1998, and 1720 MHz: Yusef-Zadeh et al. 1996).
Of the ten new masers, the two I665 MHz and four 1667 MHz main line OH masers are associated with four known OWIR stars. The two 1665 MHz masers appear along with 1667 MHz masers in the O W R stars with the brightest 1612 MHz masers. Together with the two other 1667 MHz masers, these OWIR
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Table 3 New point-like OH masers found in the 1986 VLA-BnA data set
Name
Position in B1950
Position in 52000
'usystem
Flux
(Dec.)
(R.A.)
(Dec.)
(kms-')
(mJy)
-29 03 52.9 1 -29 00 36.25
17 45 40.472 17 45 46.440
-29 05 02.27 -29 01 45.16
-29 -91
30 35
-28 -29 -29 -28
58 35.21 03 53.40 00 36.55 58 53.57
17 45 38.618 17 45 40.442 17 45 46.444 17 45 54.367
-28 -29 -29 -29
59 44.75 05 02.80 01 45.50 00 01.94
t63 -29 -82 -3
37 103 106 29
-29 -28 -28 -28
00 05.83 58 09.87 58 28.00 58 19.67
17 45 38.764 17 45 40.403 17 45 45.586 I7 45 50. 165
-29 -28 -28 -28
01 15.27
-132 +30 +39 +56
70 37 75 93
(R.A.)
New OH 1665 MHz masers OH359.880-0.087 OH359.938-0.077
17 42 29.628 17 42 35.678
New OH 1667 MHz masers 08359.952-0.036 OH359.880-0.087 08359.938-0.077 0H359.977-0.087
17 42 27.907 17 42 29.597 17 42 35.682 17 42 43.647
New OH 1720 MHz masers 0H359.930-0.049 0H359.960-0.037 OH359.966-0.056 OH359.977-0.069
17 42 28.017 17 42 29.705 17 42 34.878 17 42 39.462
59 19.16 59 37.03 59 28.20
stars are among the reddest OWIR sources with the longest periods in the GC and probably in transition to become planetary nebulae. Of the four new detections at 1720 MHz that are associated with shocks in the GC, the most interesting one is 0H359.930-0.049 at a line-of-sight velocity of - 132 km s-l. In contrast to the 1720 MHz OH masers with line-of-sight velocities between +30 to +70 km spl which are associated with the Sgr A East shell, 08359.930-0.049 and 0H359.955-0.042 (at +I32 km SKI, Yusef-Zadeh et al. 1996) originate in the CND. Their velocity and symmetry in the CND yield an enclosed mass of at least 7 . 5 lo6 ~ Mu within 47.4" ( I 3 4 pc). Granting up to 3.7 x 1 O6 MD for the black hole at the position of Sgr A* (Schodel et al. 2002), this would then imply that at least 50% of the total mass within the CND is contained in the enclosed stellar cluster and molecular cloud complexes.
References Karlsson R., Sjouwerman L.O., Sandqvist Aa., Whiteoak J.B. 2003, A&A, 403, 1011 Lindqvist M., Winnberg A,, Hahing H.J., Matthews H.E. 1992, A&AS 92,43 Sandqvist Aa., Karlsson R., Whiteoak J.B., Gardner F.F. 1987, in AIP Conf. Proc. 155, The Galactic Center, ed. D.C. Backer (AIP, New York), 95 Sandqvist Aa., Karlsson R., Whiteoak J.B. 1989, in IAU Symp. 136, the Center of thc Galaxy, ed. M. Moms (Kluwer, Dordrecht), 421 Schodel R., Ott T., Genzel R., et al. 2002, Nature 419, 694 Sjouwerman L.O., van Langevelde H.J., Winnberg A,, Habing H.J. 1998, A&AS 128, 35 Yusef-Zadeh F., Roberts D.A., Goss W.M., Frail D.A., Green A.J. 1996, ApJ 466, L25
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Astron. Nachr./AN 324. No. S 1.229 -234 (2003)/ DO1 10.1002/asna.200385I 14
Constraints on distances to Galactic Centre non-thermal filaments from HI absorption Subhashis Roy *
’
’ National Centre for Radio Astrophysics (TIFR), Pune University Campus, Post Bag No.3, Ganeshkhind, Pune 41 I 007, India. Abstract. We have studied HI absorption towards three non-thermal filaments (NTFs)Sgr C, G359.54+0.18 and G359.79+0.17 using the Giant Metrewave Radio Telescope (GMRT). Our study, for the first time, constrains the distance of the Sgr C NTF and the HI1 region seen associated with the NTF in the sky plane, to within a few hundred parsecs from the Galactic Centre (GC). A molecular cloud with a velocity of - 100 km sC1 appears to be associated with the central part of the Sgr C NTF. G359.54+0.18 shows weak HI absorption (4 CJ detection) at a velocity of - 140 !un sC1, which is the velocity of a known dense molecular cloud seen towards the NTF. This cloud is expected to be located within 200 pc from the GC and thereby provides a lower limit to the distance. The upper limit to the distance of this NTF from the Sun is 10.5 kpc. The distance to the NTF G359.79+0.17 is between 5.1 and 10.5 kpc from the Sun.
-
1 Introduction The long narrow non-thermal filaments (NTFs) observed in high resolution radio-continuum maps are 2” region of our Galaxy. These structures are less than unique features seen towards only the central 1 pc in width, but extend up to 30 pc in length. With the exception of the NTF named the Pelican (Lang et al. 1999), which is nearly parallel to the galactic plane, all other NTFs are oriented perpendicular to the Galactic plane to within 20” (Morris & Serabyn 1996, and references therein). Since these NTFs remain straight despite interaction with nearby molecular clouds, it is believed that the molecular clouds and the NTFs are in pressure equilibrium, which indicates a magnetic field strength of a few milliGauss inside the NTFs (Yusef-Zadeh & Morris 1987). Magnetic fields of comparable strengths are thought to be present in the central molecular zone (CMZ). Before any attempt is made to relate the magnetic field in the NTFs with the processes occurring in the GC, it is necessary to establish that these NTFs arc actually located in the GC region and are not chance superpositions of foreground or background objects (Lasenby et al. 1989). HI absorption towards the GC ‘Radio-arc’ (Lasenby et al. 1989) and the ‘Snake’ (Uchida et al. 1992) have indicated that they are located close to the GC, but the distances to the remaining NTFs are not constrained. Here, we present new HI absorption measurements towards three NTFs, Sgr C, G359.54+0.18 and G359.79+0.17 made with the Giant Metrewave Radio Telescope (GMRT). These observations not only constrain the distances of these objects, but also test the association of some of the above mentioned clouds with the corresponding NTFs. From single dish HI emission observations towards the NTFs under study (e.g. Cohen & Davies 1979), several large scale features are identified. Near the Galactic longitude of 359.5”, two high velocity HI emission features known as the ‘Nuclear disk’ (Rougoor & Oort 1960) and the ‘Molecular ring’ (Scoville 1972) have been found. The ‘Nuclear disk’ shows high negative velocities ranging from zz - 160 to -200 Both these features are believed km sC1,whereas, the ‘Molecular ring’ has a velocity of z - 135 km SKI.
-
* Subhashis Roy: e-rnail:
[email protected] @ 2003 WILEY-VCH Verlag GmbH C Co KGaA. Weinhem
R. Roy: HI absorption line study o f three non-thermal filaments
230
to be nearer than the GC and located at a distance of Pew hundred parsecs from it (Cohen & Davies 1979). The emission from the ‘3 kpc arm’ (Rougoor & Oort 1960) located at a distance of zz 5.1 kpc from the Sun is identified at a velocity near -53 km spl. At positive velocities, emission near 135 km spl is seen due to the HI features ‘XVI’ and ‘I’ (Cohen & Davies 1979), both of which are thought to be located behind the GC. While the feature ‘XVI’ is likely to be located within a few hundred parsecs from the GC (Cohen & Davies 1979), the feature ‘I’ is thought to be 2 kpc behind the GC (Cohen 1975). Absorption by these HI features will be used to constrain the distances to the NTFs.
2 Results In this section, we present the absorption spectra towards the target sources and identify the velocity of the HI absorption features. In all the spectra, unless stated otherwise, the X-axis represents the velocity in where, I is the observed flux density of the background km s-l and Y-axis represents the transmission (I&), source at the given frequency and 10 is the actual flux density of the source. All the velocities are expressed with respect to the Local Standard of Rest and the GC is assumed to be at a distance of 8.5 kpc.
Sgr C HI1 reeon -29
c
‘I
Nuclear disk
-65 kmls cloud
7
0
1-
e
0 ,
-200
-150
-100
-50
,
,
,Y,
, ,
0
, 50
,
,
,
,
, , 100
, ,
, 150
, ,
,
,j 200
Velocity (km sl) 1
Fig. 1 Continuum map of the Sgr C filament after high pass filtering. The Sgr C HI1 region is resolved out. The image has a resolution of 7 x 6 arcsec2, along PA=56”. RMS noise in the map is 1.6 mJy
Fig. 2 HI absorption spectrum towards the Sgr C HI1 region and the central bright part (marked ’A’ in Fig. 1) of the Sgr C NTF. The bandwidth is 2 MHz. RMS noise in the two spectra are 0.02 I and 0.034 respectively.
beam-’.
2.1
SgrC
The absorption spectra towards various parts of the object (Fig. 1) are shown in Fig. 2. Fig. 2 shows the absorption spectrum towards part ‘ A of the NTF. It shows several absorption features at negative velocities in addition to the absorption feature at 0 km s-l (line-width FZ 30 km s-’). A strong absorption feature near -54 km s-l is seen due to HI absorption by the ‘3 kpc arm’. A broad absorption feature is seen between -100 km 8 - l and -200 km s-’, with an almost linear decrease in optical depth from FZ 0.5, at -100 km s-’, to FZ 0.0 at -200 km spl. The absorption spectrum towards the Sgr C HI1 region with a resolution of 3.3 km spl show similar absorption spectrum as the NTF with a few differences, which we note here. Towards the HI1 region, the broad absorption feature seen between -100 and -200 krn s-l
Astron. Nachr./AN 324, No. SI (2003)
23 I
shows at least three main components centred at -118 km s-', -138 km S K ' and -175 km s-' with optical depth of =0.5,0.3 and 0.2 respectively. The absorption depth at these velocities are similar to what is seen towards part 'A' of the NTF.
2.2 NTF G359.79+0.17, G359.87+0. I 8 and G3S9.54+0.18 The continuum image at 20 cm of the field of NTF G359.79+0.17 and G359.54+0.18 is shown in Fig. 3 and Fig. 5 respectively. The absorption spectrum integrated over the NTF is plotted in Fig. 4 and Fig. 6 respectively. At negative velocities, absorption features can be seen near -26 km s-', and a weaker feature at -58 km s-', which coincides with the line of sight velocity of the ' 3 kpc arm'. Fig. 4 also shows the absorption spectrum towards the extragalactic source G359.87+0.18. Lazio et al. (1999) have observed HI absorption against G359.87+0.18, and the aforementioned features match with their spectrum. However, the present observations have a wider velocity coverage than Lazio et al. (1999) and we detect an additional absorption feature at + I40 km sStaguhn et al. (1998) have found a dense molecular cloud at - 140 km SK' near the bent portion of the NTF (position of the molecular cloud is denoted by 'E' in Fig. 5 ) . HI spectrum taken towards this region of the NTF (denoted by 'F' in Fig. 5 ) shows absorption (Fig. 6) at this velocity (4a detection). 50
0
100
28 57
C359.87tO.18
'\I
3-kpr arm
58
03 30 25 20 RIGHT ASCENSION (52000) Grey scale flux range= -10.9 141.7 MilliJYBEAM Con1 peak flux = 1.4175E-01 JYncr.' II"L""7 LeYP = &OWEi-03 ' (-2, -1,1,2,4,6,8,10,12, 16.20.24.32, 40.48, €4,W, 96,128) 35
17 44 40
15
P
10
Feature-1
'
Fig. 4 HI absorption spectrum towards the extragalactic source G359.87+0.18 and the NTF G359.79+0.17. The bandwidth is 4 MHz. RMS noise in the two spectra are 0.028 and 0.026 respectively.
Fig. 3 Continuum image of the NTF G359.79+0.17 ai 1.4 GHz with a resolution of 9.7 x 6.7 arcsec*, along PA=79". The rms noise is 3.0 mJy beam-'.
3 Discussion 3.1
Identification of HI features & constraints on the distances to the NTFs
Identifications of HI absorption feature is performed by comparison with features of known velocities. Absorption indicates that the continuum source is located on the far side of the HI cloud and thereby provides a constraint on the distance to the continuum source. The velocities and the distances of the known HI emission features towards the three NTFs studied here have been discussed in 4 1, which will be used to constrain the distances to the NTFs.
R. Roy: HI absorption line study of three non-thermal filaments
232
4w
-300
-200
-100
100
0
200
300
100
Velocify
Fig. 6 HI absorption spectrum integrated over the
with a resolution of 9.6 x 6.4 arcsec’, along P k 7 9 ” . The rms noise is 1.7 m l y beam-’
of the NTF where it bends (region ‘F‘in Fig. 5). The bandwidth is 4 MHz. RMS noise in the two spectra are 0.025 and 0.1 respectively. I
3.4 kpc
I1
-2 kpc
0 Sun
VeloLlly
Fig. 7 CO emission spectra towards the central part (top panel) (marked ‘A‘ in Fig. I ) and the eastern part of the NTF (middle), along with the spectrum taken towards the Sgr C HI1 region (bottom) (Data courtesy Oka et al. (1998)).
L
GC region -400 pc
8.5 kpc
Fig. 8 Schematic diagram of the Sgr C complex with HI absorbing clouds (not to scale) as seen from bottom of the Galaxy
3.1.1 SgrC Due to absorption by the line of sight HI gas and velocity crowding near 1=0”, strong absorption is observed near 0 km s-l in all spectra towards the Sgr C NTF and the HI1 region seen associated in the sky plane. Absorption by the ‘3 kpc arm’ is observed at -54 km sP1 towards the Sgr C N l T and the HI1 region (Fig. 2). The broad absorption feature (Fig. 2 ) between - I00 and -200 km s-’ is likely to be caused by several absorption features, whose line-widths are broader than their separation. The HI absorption near - 100 km sP1 is identified in CO Oka et al. (1998) emission as shown in Fig. 7. HI absorption at this velocity is caused by a cloud of size -5’, and Roy (2003) has shown it to be associated with the NTF. Absorption near - 138 km s-’ is likely to be caused by the HI associated with the ‘Molecular ring’, which has a line-width of N 40 km s-l in emission (Cohen & Davies 1979). Detection of absorption beyond - 160 km s-’ indicates absorption by the ‘Nuclear disk’ (Rougoor & Oort 1960). As ‘Molecular ring’ and the ‘Nuclear disk’ are located within 200 pc from the GC, this provides a lower limit of -8.3 kpc to these objects. We note that despite the emission feature seen in the CO map
-
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a1 -140 km s-’ (Fig. 7), no corresponding HI absorption could be detected towards the Sgr C region. The CO emission from the molecular cloud at -140 km spl is likely to be associated with the HI features ‘XVI’ and ‘I’ and both these features are located at the far side of the GC. Absence of any absorption by the HI associated with these structures indicate that the Sgr C NTF and the HI1 region are located within a few hundred parsecs at the far side of the GC, which provides an upper limit to their distances.
3.1.2 NTF G359.79+0.17, G359.87+0.18 & G359.54+0.18 No HI absorption at high positive velocity is detected towards the NTF G359.79+0.17 or G359.54+0.18. However, CO emission has been detected near +I40 km s-l towards G359.79+0.17, G359.54+0.18 and the extragalactic source G359.87+0.18, which indicates that there is no hole in feature ‘I’ (or perhaps in feature ‘XVI’) along these directions. Consequently, the upper limit to the distance of the NTF G359.79+0.17 and G359.54+0.18 is = 10.5 kpc. The presence of absorption in the spectra of the NTFs up to a negative velocity of -58 km s-l suggests absorption by the ‘3 kpc arm’ and consequently, the lower limit to their distance is N 5.1 kpc from the Sun. HI absorption near +140 km spl is only observed (Fig. 4) towards the extragalactic source G359.87+0.18. We note that both the HI emission features ‘I’ and ‘XVI’ have velocities close to 140 km s p l at this longitude and are located farther away from the GC. Therefore, the absorption near 140 km indicates that it is caused by either one or a combination of both these HI features, and is consistent with its location outside of the Galaxy. We could detect weak HI absorption at 40 level at - 140 km s-l (Fig. 6) toward part-F of the NTF shown in Fig. 5. At the same velocity, a dense molecular cloud was detected by (Staguhn et al. 1998). HI absorption by the dense molecular cloud suggests that this NTF is either embedded or located behind the CMZ found within 200 pc from the GC.
4 Conclusions HI absorption studies of three NTFs known as the Sgr C, G359.54+0.18 and G359.79+0.17 using the GMRT have yielded the following results: (a) For the first time, the Sgr C NTF and the HI1 region are shown to be located within a few hundred parsecs from the GC. (b) A molecular cloud with a velocity of -100 km s-’ appears to be associated with the central part of the Sgr C NTF. (c) HI ab5orption by the ‘3 kpc arm’ is detected against all the three NTFs, which indicates that the NTF G359.54+0.18 and G359.79+0.17 are located at a minimum distance of 5 . I kpc from the Sun. (e) Weak HI absorption (4 (7 level) at - 140 km s-l suggests that the NTF G359.54+0.18 is located at a minimum distance of = 8.5 kpc from us. (f) The maximum distance of’ the NTF G359.54+0.18 and G359.79+0.17 are estimated to be 10.5 kpc from the Sun. The present study extends the number of NTFs, which have been found to be located near the GC region to five. With most of the known NTFs now being shown near the GC, there remains little doubt that phenomena related to the central region of the Galaxy are responsible for the creation and maintenance of the NTFs.
5
Acknowledgements:
It is a pleasure to thank A. Pramesh Rao, with whom I have discussed several aspects of this work at various stages. I thank the staff of the GMRT that made these observations possible. GMRT is run by the National Centre for Radio Astrophysics of the Tata Institute of Fundamental Research.
References Cohen, R. J. 1975. MNRAS, 171,659 Cohen, R. J. & Davies, R. D. 1979, MNRAS, 186,453
234
R. Rov: HI ahsorotion line studv of three non-thermal filaments
Lang, C. C., Anantharamaiah, K. R., Kassim, N. E., & Lazio, T. J. W. 1999, ApJL, 521, L41 Lasenhy, J., Lasenhy. A. N., & Yusef-Zadeh, F. 1989, ApJ, 343, 177 Lazio, T.J. W., Anantharamaiah, K. R., Goss, W. M., Kassim, N. E., & Cordes, J. M. 1999, ApJ, 515, 196 Moms, M. & Serabyn, E. 1996, ARA&A, 34,645 Oka, T., Hasegawa, T., Sato, E, Tsuboi, M., & Miyazaki, A. 1998, ApJS, 118,455 Rougoor, G. W. 1964, Bull. Astron. Imt. Netherlands, 17, 381 Rougoor, G. W. & Oort, J. H. 1960, Proc.Nat.Acad.Sci, 46, 1 Roy, S. 2003, A&A, 403,917 Scoville, N. 2. 1972, ApJL, 175, L127 Staguhn, J., Stutzki, J., Uchida, K. I., & Yusef-Zadeh, F. 1998, A&A, 336,290 Tsuhoi, M., Handa, T., & Ukita, N. 1999, ApJS, 120, 1 Uchida, K., Morris, M., & Yusef-Zadeh, F. 1992, AJ, 104, 1533 Yusef-Zadeh, F. & Moms, M. 1987, ApJ, 322, 721
Astron. Nachr./AN 324. No. S1.235 -239 (20031 / DO1 10.1002/asna.200385034
Discovery of a non-thermal X-ray filament in the Galactic Centre Masaaki Sakano*
',', Robert S. Warwick**' ,and Anne Decourchelle****
' Department of Physics and Astronomy, University of Leicester, Leicester LEI 7RH, UK
' CEAIDSMIDAPNIA,Service d'Astrophysique, C.E. Saclay, 91 191 Gif-sur-Yvette Cedex, France Japan Society for the Promotion of Science (JSPS)
Key words The Galactic Centre, Individual: Sgr A East, Individual: XMM J174540-2904.5, X-ray We report the discovery of an X-ray filament, XMM J174540-2904.5, in the Galactic Centre region. Images from Chandru and XMM-Newton show the X-ray source is extended and coincides with a non-thermal radio structure of somewhat larger extent. The X-ray spectrum is clearly not thermal in nature, and is well approximated as a heavily absorbed power-law continuum with a photon index % 2. Combining the radio and X-ray spectra, we concluded that the emission in both wavebands probably originates in the synchrotron process. We discuss some possible origins for this peculiar non-thermal structure.
1 Introduction Many peculiar non-thermal filamentary structures have been reported on the basis of radio observations of the Galactic Centre region. Some of the filaments are relatively compact whereas others subtend large angular scales, e.g. the Galactic Centre Lobe (e.g. Mezger et al. 1996) and the jet (Sofue et al. 1989). The nature and origin of the latter features, which appear to be unique to the Galactic Centre, remains a mystery. Morris ( 1 994) summarised the properties of several individual radio filaments seen in the Galactic Centre region. Interestingly, most of them have a similar nature; they are long, thin and, although often exhibiting some curvature, are generally aligned perpendicular to the Galactic Plane. The radio emission froin these filaments is very likely to be due to the synchrotron process in which relativistic electrons spiral in a local magnetic field of -I mG. Polarisation observations identified the direction of the magnetic field to be along the filaments, i.e. orientated perpendicular to the Galactic Plane. However, recent deep VLA observations have revealed that there are many weak radio structures in the Galactic Centre region, which are more closely aligned with the Galactic Plane (e.g. Novak et al. 2003). Some of the radio non-thermal structures seen in the Galactic Centre region can probably be identified with supernova remnants, H 11 regions, or possibly pulsar nebulae. For example, Lazendic et al. (2002) detected OH maser sources from some parts of the shell of SNR G 359.1-0.5, and estimated the magnetic field to be 0.5 mG. The shell of G 359.1-0.5 is known to be interacting with the surrounding molecular cloud. Thus, the shock in the dense cloud might create the non-thermal radio structure with a strong magnetic field. Some other filaments also show morphological evidence for interactions with molecular clouds; the concentration of the matter might again enhance the particle acceleration. The new generation of X-ray telescopes carried by missions such as Chandra and XMM-Newton enable us to resolve discrete extended X-ray sources froin point sources and the overall diffuse emission in the * Corresponding author: e-mail: rnasQstar.le.ac.uk, Phone: +44 116252 3510, Fax: 4 4 116252 331 1
* * e-mail: rswOstar.le.ac.uk, Phone: +44 1162523517, Fax: +44 116252331 I '** e-mail: adecourchelleQcea.fr, Phone: +33 I690843 84. Fax: +33 169086577,
@ 2003 WILEY-VCH Verlag GmhH C Co KGaA. Wemhclm
M. Sakano et al.: A non-thermal X-ray filament in the GC
236
Galactic Centre region. Koyama (2001) has reported the discovery of a candidate X-ray filament in this region. Bamba et al. (2002) has further found some X-ray filamentary structures in the Radio Arc region. The X-ray spectra of these X-ray filaments seem featureless consistent with a non-thermal origin. Most of these X-ray structures, however, have no clear correlation with the radio filaments. The only exception is NTF G359.54+0.18, for which there is good similarity between the X-ray and radio morphology (Wang et al. 2002). The X-ray spectrum of NTF G359.54+0. I8 was, however, not determined with any precision due to poor counts statistics. In this paper, we report the XMM-Newton and Chandra discovery of a non-thermal X-ray filament, XMM 5174540-2904.5, in the Galactic Centre region, which has an apparent radio counterpart (the Sgr A-E 'wisp'=lLC 359.888-0.086=G359.88-0.07). This source is located 4 arcmin to the Galactic west of Sgr A* (see Fig. 1 Sakano et al. 2003a). Arguably, the present discovery provides the clearest example of an X-rayhadio filament in this region, both from a morphological and spectral point of view.
2 X-ray Observations We have carried out a survey with XMM-Newton concentrating on the region along the Galactic Plane within +lo of the Galactic Centre. Further details of this programme, including the mosaiced image, are presented in Warwick (2002, 2003) and Sakano et al. (2003~).XMM 5174540-2904.5 was discovered in one of the observations, designated GC6, see Sakano et a]. (2003a). We have also employed data from a Chandra observation of this region made on 8 July 2000. Further information relating to this latter observation is given in Maeda et al. (2002) and Sakano et al. (2003a).
3 X-ray and radio images
..................... . *
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Fig. 1 The Chandra image of XMM 5174540-2904.5 in the 2-8 keV band, overlaid with the 2 cm radio continuum measured by the VLA (Ho et al. 1985). The coordinates are in the 52000. This figure is taken from Sakano et al. (2003a).
Fig. 1 shows the X-ray image of XMM 5174540-2904.5 overlaid with the corresponding radio 2 cm contour. The X-ray and radio images show that the source is extended in both waveband bands and that the morphologies are well correlated in the northern region of the source (the part nearer to the Galactic
Astron. Nachr./AN 324, No. S 1 (2003)
237
Plane). This good spatial correlation strongly suggests that both the X-ray and radio cmissions have the same origin. The northern part of the source is aligned nearly perpendicular to the Galactic Plane; this nature is often seen in the radio filaments i n the Galactic Centre region, as noted earlier. On the other hand, the radio source is more extended than its X-ray counterpart and shows significant curvature at its southern extent. The radio source, which coincides with the X-ray source, was first detected in 2 and 6 cm by Ho et al. ( I 985) and designated as Sgr A-E 'wisp' based on its characteristic morphology. Lazio & Cordes ( 1 998) have later catalogued this source as ILC 359.888-0.086 based on their 1281 MHz (-20 cm) observation. They noted that this source has an angular extent of 47 arcsec, which is consistent with 2 cm and 6 cm measurements of Ho et al. (1985). More recently, Lang et al. (1999) confirmed this result with a 20 cin observation; in this case the source was designated as G359.88-0.07.
4 X-ray and radio spectra We extracted XMMIMOS and pn, and Chunclr~i/ACISspectra from an elliptical source region centred on the X-ray filament. The associated background spectra were taken from a near-sky region (see Sakano et al. 2003a for more details). If we use a thin thermal model to model the observed background-subtracted spectra, we find that a temperature in excess 40 keV is required; accordingly we rejected the thermal model. In contrast, the spectral fitting the data with a power-law model gives acceptable results. The bestfitting parameter values were: photon index r = 2.0'::; and hydrogen column density N f i = 3821, x 10""H cmp2. The observed 2-10 keV flux is 4 ~ 1 0 ~ ' ~ es-'r gcmp2, which converts to a 2-10 keV unabsorbed luminosity of 1 x 1034ergs-'. Here all the uncertainties quoted are at a 90% confidence level. Confidence contours
b)
-t Fig. 2 (a) The X-ray spectra with the best-fitting power-law model; circle (black), cross (red). star (green), and tliangle (blue) represent XMMlpn, MOS1, 2 and ChundrulACIS data, respectively. This figure is from Sakano et al. (2003a). (b) The confidence contour between the photon index (r)and the column density ( N H for ) the 6876, 90% and 99% confidence limits.
Fig. 2 shows the observed spectra together with the best-fitting power-law model and the corresponding confidence contour for column density versus photon index. The column density is found to be very large implying that the X-ray source is located in or behind dense molecular clouds. The flat power-law spectrum with heavy absorption makes this a very distinct hard source in the X-ray hardness map of the Galactic Centre region (see Fig.1 in Sakano et al. 2003b, 2003d). Combining the radio measurements at I5 GHz (2 cm), 5 GHz (6 cm) and 1.3 GHz (Ho et al. 1985; Lazio & Cordes 1998), we find all three data points to be well fitted with a power-law model with spectral index cy x 0.4 ( I E K E-"; i.e. photon index r of 1.4). Fig. 3 summarises the broad-band spectrum from the radio to the X-ray bands of the filament. We found that the measured X-ray flux sits a few decades
M. Sakano et al.: A non-thermal X-ray filament in the GC
238
below the extrapolation of the radio power-law into the X-ray band, but that the X-ray spectrum may well fit smoothly to the radio spectrum via a spectral break somewhere in the IR to EUV range. Broad band spectrum (Radio - X-ray)
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Fig. 3 The broad band energy spectrum of XMM 5174540-2904.5 from the radio to X-ray hands. The three crosses represent the radio data of 15 GHz (2 cm), 5 GHz (6 cm) (Ho et al. 1985) and 1.3 GHz (-20 cm) (Lazio & Cordes 1998), whereas the polygon represents the X-ray spectrum obtained in this work. The dashed line shows the extrapolation o f a power-law model with spectral index 0.40 (IE 0: E - Q ) which provides a best-fit to the radio data.
5 Discussion The flat X-ray spectrum of XMM J174540-2904.5 implies that the emission is non-thermal in origin. In fact, the thermal model was rejected by the spectral fitting. The broad-band spectrum from the radio to X-ray bands shows the spectral shape which can be smoothly connected somewhere between the two bands. Non-thermal X-ray emission has been detected from a number of Galactic SNRs and pulsar wind nebulae (e.g. Koyama et al. 1995). For those sources, the radio spectrum is also non-thermal and smoothly connected to the X-ray spectrum, just as in our case. The emission is most probably due to the synchrotron radiation of relativistic electrons with energy of up to -100 TeV (e.g. Reynolds & Keohane 1999). The lifetime ( T ) of synchrotron-emitting electrons is inversely proportional to the square of the magnetic field (7 c( W ' ) .In the case of the Sgr A-E 'wisp' (XMM 5174540-2904.5), Ho et al. (1985) estimated the magnetic field to be 0.3 mG, based on an equipartition argument. The lifetime of the 20 TeV electrons emitting hard X-rays through synchrotron radiation is then -7 yr. This lifetime is comparable to the observed spatial extent of the filament which is -2 light year in the X-ray band. Consequently, either the continuous injection of high-energy electrons into the source is required or, more likely, the on-going acceleration of electrons is occurring in the source itself. In fact, the position of this source coincides with the peak of the molecular cloud M-0.13-0.08 (the "20 km s-'" cloud; e.g. Mezger et al. 1986), which is only several tens of parsecs away from Sgr A' (Zylka, Mezger & Wink 1990). This together with the heavy absorption in the X-ray band strongly suggests that XMM J174540-2904.5 is embedded in, or beyond, the cloud. It is plausible that these high-density
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conditions enhance the particle acceleration hence explaining why this is a bright non-thermal source right up to the X-ray band. Coil & Ho (2000) argue from their radio observations that the radio ‘wisp’ (i.e. XMM J 174540-2904.5) may be part of a SNR (designatedas G 359.92-0.09)which lies to the side of and behind the M-0.13-0.08 cloud. If this is the case, X M M 5174540-2904.5 could be a shock front of the SNR. Interestingly, a recent Chandru deep look shows an X-ray shell-like structure, which may possibly include this XMM 5174540-2904.5 (Park et al. 2003). The X-ray colour of the opposite side of the possible X-ray shell is, however, very different from XMM J 174540-2904.5. Another bright radio structure Sgr A-D designated by Ho et al. ( 1985), which is regarded as a part of their proposed SNR, has no apparent X-ray counterpart, with only XMM J 174540-2904.5 being bright in the hard X-ray band. Consequently, the interpretation of the X-ray filament in terms of a SNR shell remains uncertain. An alternative is that we are dealing with an isolated X - r a y h d i o filament. The alignment of the X-ray source, perpendicular to the Galactic Plane, is then consistent with the phenomenology of other isolated radio non-thermal filament in the Galactic Centre region. Or this source could possibly be an extragalactic background object, namely a one-sided jet emanating from a QSO, although the lack of a central point source corresponding to the QSO core somewhat weakens this argument. In any case, this is the first clear detection of an X-ray filament which has a radio counterpart and unequivocally has a non-thermal X-ray spectrum. Further detailed investigation in the X-rayhadio bands, particularly via deep high resolution imaging observations, should help reveal the true nature of this peculiar source. Acknowledgements MS acknowledges the financial support from JSPS.
References Bamba, A., Murakami. H., Senda, A., Takagi, S., Yokogawa, J., & Koyama, K. 2002, Proc. New Visions of the X-ray Universe in the XMM-Newton and Chandra era, in press (astro-ph/0202010) Coil, A. L., & Ho, P. T. P. 2000, ApJ, 533, 245 Ho, P. T. P., Jackson, J. M., Barrett, A. H., & Armstrong, J. T. 1985, ApJ, 288, 575 Koyama, K. 2001, in H. Inoue & H. Kunieda, ed., ASP Conf. Ser. Vol. 251, New Century of X-ray Astronomy, Astron. SOC.Pac., San Francisco, p.50 Koyama, K., Petre, R., Gotthelf. E. V., Hwang, U., Matsurd, M., Ozaki, M., & Holt, S. S. 1995, Nature, 378, 255 Lang, C. C., Moms, M., & Echevarria, L. 1999, ApJ, 526,727 Lazendic, J. S., et al. 2002, MNRAS, 331, 537L Lazio, T. J. W., & Cordes, J. M. 1998, ApJS. 118, 201 Maeda, Y., et al. 2002, ApJ, 570,671 Mezger, P. G., Chini, R., Kreysa, E., & GemUnd, H. -P. 1986, A&A, 160,324 Mezger, P. G., Duschl, W. J., & Zylka, R. 1996, A&AR, 7,289 Morris, M., 1994, in R. Genzel& A.1. Harris, eds., The nuclei of Normal Galaxies. Kluwer, Dordrecht, p.185 Novak, G., et al. 2003, these proceedings Park, S., et al. 2003, these proceedings Reynolds, S. P., & Keohane, J. W. 1999, ApJ, 525, 368 Sakano, M., Warwick, R. S., Decourchelle, A., & Predehl, P. 2003a, MNRAS, 340, 747 Sakano, M., Warwick, R. S., & Decourchelle, A. 2003b, AdSpR, submitted Sakano, M., Warwick, R. S.,& Decourchelle, A. 2003c, Proc. “Japan-GermanyWorkshop on Galaxies and Clusters of Galaxies”. p.9 (astro-ph/0212464) Sakano, M., Warwick, R. S., & Decourchelle, A. 2003d, these proceedings Sofue, Y.,Reich, W., &L Reich, P. 1989, ApJL, 341, L47 Wang, Q. D., Gotthelf, E. V., & Lang, C. C. 2002, Nature, 415, 148 Warwick, R. S. 2002, Proc. New Visions of the X-ray Universe in the XMM-Newton and Chandra era, in press (astro-pW0203333) Warwick, R. S. 2003, these proceedings Zylka, R., Mezger, P. G., & Wink, J. E. 1990, A&A, 234, 133
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Astron. Nachr./AN 324, No. S 1,24 1-245 (2003) / DO1 10.1002/asna.200385035
High-negative velocities in the inner 25 pc of the Galactic center Lorht 0. Sjouwerman National Radio Astronomy Observatory, P.0. Box 0, Socorro, NM 87801; Isjouwerman@nrao. edu
Key words Galactic center, high-velocity stars, high-velocity gas, kinematics
Almost three decades ago, in a survey for OH masers in the Galactic center, Baud et al. detected an OH/IR star with a line-of-sight velocity of -343 !an sK1 with respect to the Local Standard of Rest, 50 pc in projection from Sgr A*. Since then, at least three more high-velocity (IVl > 250 krn s-') O M R stars have been found within about 25 projected pc of the Galactic center, all with a blue-shifted (negative) velocity. Over the years, several authors have also found emission and absorption by high-negative velocity gas (HI,HzCO, HCO', CO, OH) toward the Galactic center; gas with IVI > 180 km sK1.This contribution attempts to explain the observations of the high-negative velocity OWIR stars by relating their remarkable spatial and kinematic alignment in the inner 25 pc to the OhSerVdtiOnS of the high-negative velocity gas.
1 Introduction: stars on highly elongated orbits While performing a survey for OH sources in the direction of the Galactic center (GC), Baud et al. (1975) observed a double peaked 1612 MHz OH absorption feature in a frequency switched observation, 25' from Sgr A* (50 pc in projection). Further analysis showed that this absorption feature was actually a double peaked emission feature in the reference band. As the double peaked profile is characteristic for OWIR stars, this star therefore had to be at a high velocity, with a velocity of about -340 km s-' (all velocities in this paper refer to line-of-sight velocities with respect to the Local Standard of Rest). Although it was known that high velocities (up to f 2 0 0 km S C ' ) are not uncommon in the dircction of the GC, the authors already remarked that the star is far outside that range. Another striking characteristic is that the object seems to be moving counter to the Galactic rotation. In subsequent papers where they analyzed their results, it was demonstrated that high-velocity stars, i.e. stars with / V / > 250 km s-', should be uncommon (Baud et al. 1979), and that a (high) negative velocity due to a star outside thc Solar circle was unlikely to explain the observation (Baud et al. 1981). About ten years later, in 1992, van Langevelde et al. published the detection of two additional highvelocity OH/IR stars. They are located even closer to the GC, within 12' or within 25 pc in projection, and also appear at negative line-of-sight velocities of about -300 to -350 km s- l ; both moving with Galactic rotation. No OWIR stars were found at high-positive velocities. Having three high-negative velocity (HNV) OH/IR stars among about 150 (from a survey by Lindqvist et al. 1992), the authors were able to distinguish between some possible explanations. As the possibility of finding three HNV stars with the same sign is 25%, thcy did not consider this as important. However, they noted that 2 3 HNV stars is a few percent, which is too high to be explained by a high-velocity tail of an isotropic velocity distribution, and that collision mechanisms also would produce velocities higher than the apparent range up to the roughly 350 km s K 1 observed. They concluded - conditional upon no significant additional detections of negative high-velocity stars without corresponding detections of positive high-velocity stars - that high-velocity stars were consistent with low angular momentum, highly elongated orbits along the line of sight in a tri-axial bulge of bound bulge stars currently passing the GC. 0Lo03 WlLEY-VCH
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2 High-velocity gas on hyperbolic orbits Two decades ago, Gusten & Downes (1981) pointed out that, toward the radio continuum of the GC, absorption by H2C0 and H I gas was seen at a velocity around - 190 km s-’. Furthermore they remarked that this velocity exceeds the velocity range of the known gaseous components, that the gas must be in front of the GC, and that equilibrium rotation about the nucleus could be ruled out (unless the center of rotation was shifted 210’ away from the nucleus). They favored an “ejection” model after they argued against a large scale foreground “field” model. About ten years later, Man et al. (1992) suggested that some of this HNV gas interacts with the GC environment, whereas Pauls et al. (1993) proposed the HNV gas close to the GC to be falling inward from behind. In 1993, Liszt & Burton discussed the observation of Marr et al. (1 992) and Pads et a1 (1 993) that Sgr A * was not found to be occulted by the HNV gas (see however Karlsson et al. 2003a, 2003b). and argued that the HNV gas was a foreground cloud and unrelated to the phenomena in the GC. They mainly based this on their CO emission mapping of the HNV gas, which showed an elongated East-West structure West of Sgr A* with a very small velocity dispersion across the feature, and which is suspiciously parallel to the rotation axis of the Circumnuclear Disk (CND), as well as to the emission of the near side kpc-scale “expanding” molecular gas. Many papers followed with different tracers and on different angular scales in which the physical association between the HNV gas and Sgr A* with its direct environment was investigated (e.g. Yusef-Zadeh et al. 1993; Marshall & Lasenby 1994; Yusef-Zadeh et al. 1995; Liszt & Burton 1995; Zhao et al. 1995; Roberts et al. 1996; Sofue 1996; Karlsson et al. 2003a, 2003b). Arguments in “favor of this”, and arguing “against that” went back and forth, and it probably now would be accepted that the HNV gas and the GC do interact, in particular at a location a few arcminutes to the South-West of Sgr A*. The HNV gas has been modeled as gas on a tilted hyperbolic orbit about Sgr A*, where the gas is tidally disrupted within the central 100 pc and is flowing toward the nucleus from the far side (Zhao et al., 1995; Roberts et al., 1996). However, none of the authors directly related the observations of the HNV stars and the HNV gas.
3 Associating the HNV stars and HNV gas: the emerging picture In a deep OH survey, Sjouwerman et al. (1998; Sjouwerman 1997) found another high-velocity OWIR star, again at a negative velocity of -301 km s-l. Actually, the new HNV O M R star appears almost exactly in the line connecting the two HNV stars of van Langevelde et al. (1992) in the sky plane. This alignment still holds when another HNV star at -346 km spl is included, a HNV star reported by Menten & Reid (1998, in their figure and recently associated with IRS 9 by Reid et al. 2003). Finding significantly more, negative-only high-velocity stars in the velocity range -300 to -350 km spl clearly undermines the tri-axial bulge, bound elongated orbit HNV explanation by van Langevelde et al. (1992) as high-positive velocities have been equally well surveyed in these studies. Investigating the striking East-West orientation of the spatial alignment of the HNV stars, in the same direction as the CO-cloud of Liszt & Burton (1993), Sjouwerman (1997) linked the observations of the HNV stars with the observed HNV gas (Fig. I). It turned out that the large scale HNV gas distribution, as well as the HNV stars, lie in a plane with the same inclination with respect to the Galactic plane, which does not intersect with the position of Sgr A* itself (but South-West of it). Moreover, the velocity structure of the HNV gas also aligns with the velocities of the HNV stars in the longitude-velocity diagram (Fig. 2). Recalling the explanation of the HNV gas with an inclined hyperbolic orbit by Zhao et al. (1995) and Roberts et al. (1996), it seems plausible to investigate explaining the HNV stars with this model. Picturing that the gas and stars are falling in from a positive longitude somewhere (North-)East and behind the GC. on an inclined orbit with the peribothron (periastron) South-West of Sgr A* and elongated along the line of sight, the gas is gravitationally more easily influenced by the GC than the stars. The gas interacts, decelerates, and dissipates, giving rise to a hyperbolic appearing orbit (a transition of xl-to-xz orbits?), and the observedinteractions of the HNV gas in the inner few pc. The unbound heavy stars, the HNV stars with
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g a l a c t i c longitude o f f s e t ( a r c m i n ) Fig. 1 Galactic sky distribution of HNV OH/IR stars and HNV gas (here the "CO gas; Liszt & Burton 1993). Offsets are with respect to Sgr A*. The arbitrary contours (gray-scales)outline the HNV gas with a velocity of - 190 (-213) !an sP1. The HNV stars at a velocity of around -320 (-150) h s-l are the open (closed) symbols. Comparison with other HNV gas, for example H I absorption (Yusef-Zadehet al. 1993) or 1667 MHz OH absorption (Karlsson et al., 2003b, and these proceedings), shows a similar picture. Taken from Sjouwennan (1997).
-360 < V < -290 km s-', are hardly affected, nor decelerated, and continue almost in a straight line, i.e. more to the South (more negative longitudes) of the HNV gas. The discussion of IRS 9 (at V= -346 kin s-l) being unbound by Reid et al. (2003) would support this vision. It would explain the high velocities of the HNV stars, their location with respect to the GC and HNV gas, the bound stars at -150 km s-' (that may have decelerated and thus follow the HNV gas closer), and the asymmetry of lacking positive high-velocity components. That the gas and stars have a similar inclination as the rotation axis of (but offset from) the CND (Liszt & Burton 1993) and the kpc-scale radio lobe (Sofue 1996) should be regarded
Lorint 0. Sjouwerman: High-negative velocities in the inner 25 pc of the Galactic center
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Galactic longitude ~ f f f s e t(arcrnin) Fig. 2 Galactic velocity distribution of OH/IR stars (Sjouwerman et a]. 1998). The HNV OWIR stars with velocities around -320 km s-l clearly stand out (open star symbols in Fig. I), as well as the group around -150 km sC1 and with a Galactic longitude between 35' of Sgr A" (filled star symbols in Fig. 1). The circles are the OWIR stars known prior to the Sjouwennan et al. (1998) survey. The spatial and velocity structure of the HNV gas nicely fits into the diagram between the -320 km s-' and - IS0 km s-' groups. Modified from Sjouwennan (1997).
as coincidental, although the 1720 MHz OH masers at +60 km s-' (A, D, E, F, and G, Yusef-Zadeh et al., 1996, 1999) that also align might hypothetically result from the deflected gas.
4 Remains to be explained: Baud's star ! Ironically, this model does not directly solve the high-velocity problem posed by Baud's star, the star that triggered one to start thinking about this problem three decades ago. Possible solutions, that all can explain
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the retrograde circular motion and the deviation from the straight alignment with the H N V gas and other HNV stars. are that Baud's star: resulted from a tolally unrelated single event, e.g. such a s the collision mechanism described by van Langevelde et al. ( I 992).
falls in from the other side, indeed having an opposite orbital angular momentum. This would agree with the velocity of the other H N V stars, but one would expect a better alignment in the plane of the sky.
lags in a similar orbit, still o n the other sidc of S g r A*. This would explain the difficulty of finding a red counterpart (Blommaert e t al. 1998) and is consistent with a peribothron (periastron) South of Sgr
A*. g o t deflected by Sgr A*, but still moves almost in the line of sight. This would agree with the misalignment (to the proper side of Sgr A*) with the other H N V stars, and may b e indicated by the HNV gas East of Sgr A* as observed by Yusef-Zadeh et al. (1993, their Fig. 2). It is also consisten1 with the proper motion vector of the probably unbound H N V source IRS 9 (Reid e t al. 2003).
5
Summary
The explanation of the high-velocity stars by van Langevelde et al. (1992) does not hold because significantly more negative high-velocity stars have been found. O n the other hand, the high-velocity stars d o align with the high-negative velocity gas that i s intcracting with the Galactic center. Modeling of the gas (Zhao et al. 1995;Roberts et al. 1996),indications that the high-negativevelocity stars are not bound (Reid e t al. 2003), and combination of the kinematic observations of the gas and stars, picture a cloud falling in from the far side of, and interacting with the Galactic center. T h e unbound high-velocity stars were filtered from the cloud by the interaction, where Baud's star was possibly deflected by S g r A*.
References Baud, B., Habing, H. J . , Osullivdn, J. D., Winnberg, A. & Matthews, H. E. 1975, Nature 258,406 Baud B, Habing, H. J., Matrhews, H. E. & Winnbeg, A. 1979, A&AS 35, 179 Baud B, Habing, H. J., Matthews, H. E. & Winnberg, A . 1981, A&A 95, 171 Blommaert, J. A. D. L., van der Veen, W. E. C. J., van Langevelde, H. J . , Habing, H. J. & Sjouwerman, L. 0. 1998, A&A 339,991 Gusten, R. & Downes, D. 1981, A&A 99, 27 Karlsson, R., Sandqvist, Aa., Sjouwerman, L. 0. & Whiteoak JB 2003a, these proceedings, Sect. 111 Karlsson, R., Sjouwerman, L. 0..Sandqvist, Aa. & Whiteoak JB 2003b, accepted by A&A Lindqvist M, Winnberg, A., Habing, H. J. & Matthews, H. E. 1992, A&AS 92, 43 Liwt, H. S. & Burton, W. B. 1993, ApJ 407, L25 Liszt, H. S. & Burton, W. B. 1995, ApJS 98, 679 Marr, J. M.. Rudolph, A. L., Pauls, T. A., Wright, M. C. H., & Backer, D. C. 1992, ApJ 400, L29 Marshall, J. & Lasenby, A. 1994, MNRAS 269, 619 Menten. K. M. & Reid, M. J . 1998, ASP Conf. Scr. 144 IAU Colloq. 164, 229 Pauls, T., Johnston, K. J., Wilson, T. L., Man-, J. M. & Rudolph, A. 1993. ApJ 403, L13 Reid, M. J . , Menten, K. M., Genzel. R., Ott, T., Schodel, R. & Eckart, A. 2003, ApJ in press, (astro-ph 0212273) Roberts, D. A , , Yusef-Zddeh, F. & Goss, W. M. 1996, ApJ 459, 627 Sjouwerman, L. 0. 1997, Onsala PhD thesis, Chalmers University of Technology, Gothenburg, Swedcn Sjouwerman, L. 0..van Langevelde, H. J., Winnberg, A. & Habing, H. J. 1998, A&AS 128, 35 Sofue, Y, 1996, ApJ 459, L69 van Langevelde, H. J., Brown, A. G. A,, Lindqvist, M., Habing, H. J. & de Zeeuw, P. T. 1992, ApJ 396, 686 Yusef-Zadeh, F., Lasenby, A. &Marshall, J. 1993, ApJ 410, L27 Yusef-Zadeh, F., Zhao, J.-H. & Goss, W. M. 1995, ApJ 442, 646 Yusef-Zddeh, F., Roberts, D. A,, Goss, W. M., Frail, D. A. & Green, A. J. 1996, ApJ 466, L25 Yusef-Zadeh, F., Roberts, D. A,, Goss, W. M., Frail, D. A. & Green, A. 5. 1999, ApJ 512, 230 Zhao, J.-H., Goss, W. M. & Ho, P. T. P. 1995, ApJ 450. 122
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Astron. Nachr./AN 324, No. S I , 247 - 253 (2003)/ DO1 10.1002/asna.200385054
Really Cool Stars and the Star Formation History at the Galactic Center
',
Robert D. Blum" Solange V. Ramirez**I,Kristen Sellgren***3,and Knut Olsent I
*
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Cerro Tololo Interamerican Observatory, Cassila 603, La Serena, Chile SIRTF Science Center, JPL/Caltech, Pasadena. CA 91 125, USA Astronomy Department, The Ohio State University, 140 West 18th Ave, Columbus, OH 43210, USA
Key words Galactic Center, AGB Stars, Star Formation Abstract. We present AlAA = 550 to 1200 near infrared I1 and K spectra for a magnitude limited sample of 79 asymptotic giant branch and cool supergiant stars in the central h 5 pc (diameter) of the Galaxy. Using a set of similar spectra ohtained for solar neighborhood stars with the same range i n T..Rand MI, M > 1 Ma are taken to have their remnant mass at the present time, then the present inass in stars is reduced by about 45 ‘ro (i.e. we infer the present mass in stars remaining in the cluster to be 0.55 x the total mass formed over all times). The mass N
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Fig. 3 Hertzprung-Russell diagram for the Galactic Center (GC) stars (shown as open circles with typical uncertainty given by the error bars in the upper left comer of each panel) and comparison stars (Comp Stars, plotted as filled diamonds). The GC stars analyzed at high spectral resolution by Ramhez et al. (2000) and Carr et al. (2000) are plotted as filled squares. Model isochrones are shown for reference. The isochrones are from Bertelli et al. (1994) for age < 100 Myr and Girardi et al. (2000) otherwise. Isochrones are plotted for ages of 10 Myr, 25 Myr, 50 Myr, 100 Myr, 500 Myr, 1 Gyr, 4 Gyr, and 12 Gyr. The models in the lefl panel have [Fern] = 0.0, and the models in the right panel have [Fern] = -0.2. Neither set of isochrones reaches the coolest stars, but the [Fe/H] = 0.0 isochrones extend to cooler temperatures and thus fit more Galactic stars than the [Fern] = -0.2 isochrones. Comparison to the isochrones shows that all the GC stars classified as giants are AGB stars; they are too luminous to be first ascent giants. This is a consequence of the selection criteria. The horizontal line segment at MI,,,^ = -7.2 in each panel indicates the approximate observed luminosity above which only supergiants lie BSD96.
lost through stellar winds could be expelled from the region andlor recycled into new generations of stars. The total present mass in stars and remnants due to the cumulative star formation history inside the central 2 pc (radius) is thus 56.6 f 2 x lo6 M, and so within about a factor of two compared to expectations from the dynamical models. Changes in the lower mass cut-off or slope in the adopted IMF could result in less total star formation. For example, simply cutting off the mass function at one Ma would reduce the total by approximately a factor of 2. Finally, we have implicitly assumed all the tracers of the SFH lie within a true radius of 2 pc, but they are actually distributed in a projected radius of 2 pc. Given the steepness of the stellar cluster radial density distribution (20.5 pc core radius), the overestimate is likely to be small. A full discussion and updated results for this work are given by Blum et al. (2003)
References Bertelli, G.. Bressan, A,, Chiosi, C., Fagotto, F., & Nasi, E. 1994, A&AS, 106, 275 Binney, J . , Gerhard, 0. E., Stark, A. A,, Bally, J., & Uchida, K. 1. 1991, MNRAS, 252,210 Blum, R. D., Sellgren, K. &DePoy, D. L. 1996, ApJ, 470,864
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Model A
Age (Gyr)
4
Age (Gyr)
b)
Fig. 4 Star formation history (SFH) for the Galactic center. The crosses represent the results to the SFH fits to the Hertzspmng-Russell Diagram (Figure 3) for Models A (upper panel) with Solar [Fe/H] throughout and Model B (lower panel) with [Fe/H] = -0.2 in the oldest age bin. Model A is a statistically better fit; see text. The age bins corresponding to the horizontal width of the crosses are 10 Myr - 100 Myr, 100 Myr - 1 Gyr, 1 Gyr - 5 Gyr, and 5 Gyr - 12 Gyr. The vertical height of each cross is the onc sigma error in the star formation rate for the respective bin; see text. Blum, R. D., Sellgren, K. &DePoy, D. L. 1996, AJ, 112, 1988 (BSD96) Blum R. D., Ramirez, S. V., Sellgren, K., & Olscn, K. 2003, ApJ, in press Blitz, L. & Spergel, D. N. 1991, ApJ, 379. 631 Cam, J. S., Sellgren, K., & Balachandran, S. C. 2000, ApJ, 530, 307 DePoy, D. L., Gregory, B., Elias, J., Montane, A., Perez, G., & Smith, R. 1990, PASP, 102, 1433 DePoy, D. L., Atwood, B., Byard, P., Frogel, J . A., & O'Brien, T. 1993, in SPIE 1946, "Infrared Detectors and Instrumentation," pg 667 Dolphin, A. E. 2002, astro-pW0112331 Dwek. E. et al. 1995, ApJ, 445,716 Dyck, H. M., Benson, J. A., van Belle, G. T., & Ridgway, S. T. 1996, AJ, 11 I , 1705 Dyck, H. M., van Belle, G. T., & Thompson, R. R. 1998, AJ, I 16, 981 Frogel, J. A., Persson, S. E., Mdtthews, K., & Aaronson, M. 1978, ApJ, 220. 75 Frogel, J. A. ti Whitford, A. E. 1987, ApJ, 320. 199 Genzel, R., Hollenbach, D., & Townes, C. H. 1994, Reports of Progress in Physics, 57,417 Ghez, A. M., Klein, B. L., Moms, M., & Becklin, E. E. 1998, ApJ, 509, 678 Girardi, L., Bressan, A., Bertelli, G., & Chiosi, C. 2000, A&AS, 141, 371 Giisten, R. 1989, IAU Symp. 136: The Center of the Galaxy, 136,89 Giisten, R., Genzel, R., Wright, M. C. H., Jaffe. D. T., Stutzki, J., & Harris, A. 1. 1987, ApJ, 318, 124 Jackson, J. M., Geis, N., Genzel, R., Hams, A. I., Madden, S., Poglitsch, A., Stacey, G. J., & Townes, C. H. 1993, ApJ, 402, 173 Kleinmann, S. G. & Hall, D. N. B. 1986, ApJS, 62, 501 Liszt, H. S. & Burton, W. B. 1980, ApJ, 236, 779 Mezger, P. G., Zylka, R., Philipp, S., & Launhardt, R. 1999, A&A, 348,457 Miralda-EscudC, J. & Gould, A. 2000, ApJ, 545, 847 Moms, M. & Serabyn, E. 1996, ARAA, 34,645 m186 Mulder, W. A. & Liem, B. T. 1986, A&A. 157, 148 Olsen, K. A. G. 1Y9Y, AJ, 1 17. 2244 Pogge, R. W. et al. 1998, SPIE, 3354,414 Ramirez, S. V., Sellgren, K., Carr, J. S., Balachandran, S. C., Blum. R., Terndrup, D. M., & Steed, A. 2000, ApJ, 537,205 Salpeter, E. E. 1955, ApJ, 121, 161 Sanders, R. H. 1999, ASP Conf. Ser. 186: Thc Central Parsecs of the Galaxy, 250 Weiland, J. L. el al. 1994, ApJ , 425, L81
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Astron. Nachr./AN 324. No. S 1.255 -261 (2003) / DO1 10.1002/asna.200385060
Massive Stars and The Creation of our Galactic Center Donald F. Figer'
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' STScI, 3700 San Martin Drive, Baltimore, MD 21218 Key words Galactic Center, Stellar Clusters, Star Formation Abstract. Our Galactic Center hosts over 10% of the known massive stars in the Galaxy, the majority of which are located in three particularly massive clusters that formed within the past 5 Myr. While these clusters are extraordinary, their formation repesents about half of the total inferred star formation rate in the Galactic Center. There is mounting evidence that the clusters are just present-day examples of the hundreds of such similar clusters that must have been created in the past, and whose stars now comprise the bulk of all stars seen in the region. I discuss the massive stellar content in the Galactic Center and present new data obtained with HST/NICMOS and Geniini AO, and an analysis that suggests that effects of continuous star formation in the Galactic Center can be wen in the observed luminosity functions.
1 Introduction Over 10% of the known massive stars (Mi,,it>20 M o ) in the Galaxy reside in three clusters of young stars located within 30 pc of the Galactic Center. These clusters are the most massive young clusters in the Galaxy and contain approximately 30 Wolf-Rayet (WR) stars, at least 2 Luminous Blue Variables (LBV), approximately a half dozen red supergiants, and approximately 450 0 stars. Together, they emit enough ionizing radiation to account for roughly half of the thermal radio emission in the central few degrees of the Galaxy, suggesting that the young clusters contain approximately half of the stars recently formed in this region. An additional collection of young stars exists in the region, with members scattered about the central 50 pc; some have evolved to the WR stage, while others are still deeply embedded within their natal dust cocoons. A lower bound to the current star formation rate can be approximated by dividing the mass in the clusters by their ages, i.e. 5(104) M 0 / 5 M y r ~ O . 0 1Mcj/yr, or a star formation rate density of Mo/yr p c 3 . This rate is approximately 250 times higher than the mean rate in the Galaxy, and about the same factor lower than the rate in starhurst galaxies. Clearly, the Galactic Center has formed a plethora of stars in the past 5 Myr, hut it is less apparent when the nlillions of stars in the central 50 pc formed. If we assume that the star formation rate in the past was similar to the present rate, then the total mass of stars formed over the past 10 Gyr is -1 O8 Ma within a radius of 30 pc of the Galactic Center, and an order of magnitude greater than this amount over the whole Central Molecular Zone, as first suggested by Serabyn & Morris ( I 996).
2 The Central Cluster The Central Cluster contains over 30 evolved massive stars having Minitial >20 Mo (Becklin & Neugehauer 1968, Rieke & Lehofsky 1982, Lebofsky, Rieke, & Tokunaga 1983, Forrest et al. 1987, Allen, Hyland, & Hillier 1990, Krahbe et al. 1991; 1995, Allen 1994, Rieke & Rieke 1994, Blum et al. 1995, Eckart et al. 1995, Genzel et al. 1996, Tamhlyn et al. 1996). A current estimate of the young population * Corresponding author: e-mail: figerQstsci.edu, Phone: 410-453-9321 Fax: 410-516-2829
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includes 9 WR stars, 20 stars with OfpelWN9-like K-band spectra, several red supergiants, and many luminous mid-infrared sources in a region of 1.6 pc in diameter centered on Sgr A* (Genzel et al. 1996). In addition, I estimate that it contains at least 100 0-stars (07 and later) still on the main sequence. Najarro et al. (1994) and Najarro et al. (1997) modeled the infrared spectrum of 8 blue supergiants in the center, finding characteristics consistent with an “OfpeIWN9” classification, and concluding that they formed a few million years ago. More recently, Paumard et al. (2001) reviewed the emission-line stellar population in the central parsec, using new narrow-hand infrared imaging. Eckart et al. (1999) and Figer et al. (2000) identified massives stars within a few thousand AU of the supennassive black hole. Evidently, a significant fraction of this small group of stars are young ( T lo1’ their present location would suggest. Results from proper-motion studies suggest that at least some of the stars in this cluster are bound to the black hole and are not on highly elliptical orbits (Ghez et al. 2001; Eckart et al. 2002); therefore they are likely to he near to their formation sites.
3 The Quintuplet Cluster The Quintuplet Cluster is located approximately 30 pc, in projection, to the northeast of the Galactic Center (Glass, Catchpole, & Whitelock 1987). In addition to the five bright stars for which the Quintuplet was named (Nagata et al. 1990; Okuda et al. 1990),the Quintuplet cluster contains a variety of massive stars, including four WN, five WC (possibly ten, see below), two WN9/0fpe, two LBV, one red supergiant and several dozen less-evolved blue supergiants (Figer et al. 1999a, 1999~).The five Quintuplet-proper members are massive stars (L lo5 L o ) embedded within dusty cocoons, although their spectral types and evolutionary status are unknown (Moneti et al. 2001). Figer et al. (1996, 1999a) argue that these stars are dust-enshrouded WCL stars, similar to WR 140 (Monnier et al. 2002) and WR 98A (Monnier et al. 1999). New evidence in support of this hypothesis was presented by Chiar (2003), Law (2003), and Lang (2003). In addition to these post-main sequence stars, it is likely that 100 0-stars still on the main sequence exist in the cluster, assuming a flat to Salpeter IMF. The total cluster mass is estimated to he w 1 O 4 M a . The total ionizing flux is photons spl, enough to ionize the nearby “Sickle” HI1 region ((30.18-0.04). The total luminosity from the massive cluster stars is NN 107.5La,enough to account for the heating of the nearby molecular cloud, M0.20-0.033. The two LBVs in the cluster are added to the list of 6 LBVs in the Galaxy. They include the Pistol Star (Moneti et al. 1994; Figer et al. 1995a, 1995b, 1998, 1999b; Cotera 1995; Cotera et al. 1996), one of the most luminous stars known, and a newly identified LBV (Geballe et al. 2000) that is nearly as luminous as the Pistol Star. Most of the luminous stars in the cluster are thought to be 3-5 Myr old. N
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4 The Arches Cluster First discovered about 10 years ago as a compact collection of a dozen or so emission-line stars (Cotera et al. 1992; Nagata et al. 1995; Figer 1995a; Cotera 1995; Cotera et al. 1996; Blum et al. 2001), the Arches cluster contains thousands of stars, including at least 160 0 stars, according to Figer et al. (1999~).Figer et al. (1999~) used HST/NICMOS observations to estimate a total cluster mass (2104M a ) and radius (0.2 pc) to arrive at an average mass density of 3(105) M a pcP3 in stars, suggesting that the Arches cluster is the densest, and one of the most massive, young clusters in the Galaxy. They further used these data to estimate an initial mass function (IMF) which is relatively flat (r -0.61t0.1) with respect to what has been found -1.35, Salpeter 1955) and other Galactic clusters (Scalo 1998). Stoke et for the solar neighborhood (I? al. (2002) recently confirmed this flat slope by analyzing the same data and recently obtained Gemini A 0 data. Figer et al. (2002) estimated an age of 2.510.5 Myr, based on the magnitudes, colors, mix of spectral
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Fig. 1 (left) K-band spectra of four Wolf-Rayet stars (WNL) in the Arches Cluster from Figer et al. (2002). Emission lines can be seen at 2.058 p m (HeI), 2.104 pm (NIII), 2.1 12/113 p m (HeI), 2.1 15 pm (NIII), 2.166 pm (HeVHI), 2.189 pm (HeII), and 2.224/225 ,urn (NIII). The sharp feature near 2.32 pm in the spectrum for star #2 is due to a detector defect, and the absorption features longward of 2.33 pm are due to imperfect correction for telluric absorption. (right) Equivalent width of the 1.87 pm feature in massive stars of the Arches Cluster (Figer et al. 2002b; see also Blum et at. 2001) as a function of apparent magnitude in the HST/NICMOS F205W filter. The feature has contributions from He1 and HI.
10 5 0 - 5 - 1 0 R.A. o f f s e t {arcseconds} Fig. 2 Difference image, F187N-FI90N. highlighting stars with excess emission at 1.87 pm. Radio sources (v 4.9 G H z ) are shown by circles, as identified in Lang et al. (2001), and as squares, as identified by Figer et al. (2002). X-ray sources are shown by diamonds, as identified by Yusef-Zadeh et al. (2002).
types, and quantitative spectral analysis of stars in the cluster. Given the current state of knowledge about this cluster, it now seems apparent that we have observed a firm upper-mass cutoff (Figer 2003a). Note Indeed, we should even expect that we should expect at least 10 stars more massive than M,nitlal=300Ma. one star with an initial mass of 1,000 Mo ! Of course, it is questionable how long such a star would live; however, it is clear that the Arches cluster IMF cuts off at around 150 Ma.Finally, even if we steepen thc IMF slope to the Salpeter value, we still should expect at least 4 stars more massive than 300 Mo. Figer et al. (2002) conclude that the most massive stars are bona-fide Wolf-Rayet (WR) stars and are some of the most massive stars known, having Mirritial> 100Ma, and prodigious winds, A? > lop5 M a y-',
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that are enriched with helium and nitrogen. These findings are largely based upon the spectra and narrowband equivalent widths shown in Figure 1, and a detailed quantitative analysis of these data (also see Najarro 2002). Figer et al. (2002) found an upper limit to the velocity dispersion of 22 km s-', implying an upper limit to the cluster mass of 7(104) M a within a radius of 0.23 pc, and a bulk velocity of vcluster xi-55 km s-' for the cluster. It appears that the cluster happens to be ionizing, and approaching, the surface of a background molecular cloud, thus producing the Thermal Arched Filaments. They estimate that the cluster produces 4(1051) ionizing photons s-', more than enough to account for the observed thermal radio flux from the nearby cloud. Commensurately, it produces lo7.' La in total luminosity, providing the heating source for the nearby molecular cloud, L,-loud = lo7 Lo. These interactions between a cluster of hot stars and a wayward molecular cloud are similar to those seen in the "Quintuplet/Sickle" region. Finally, note that significant work is being done on this cluster at radio and x-ray wavelengths, i.e. shown in Figure 2.
5 The Star Formation History of the Galactic Center The evidence for recent (-- 30 mag. For DB00-4,5,6 and 58, the agreement between the X-ray and IR derived values is good. Av values for these two regions suggest that they are outside of the GC. However, DBO 1-42 shows a significant difference between the X-ray and IR derived Av . Upon closer inspection, the X-ray/IR correlation is relatively weak, since the PSF is relatively large in this case. Aside from this possible false X-ray/IR correlation, we arc seeing highly absorbed X-ray emission from the candidate cluster. The color-color diagram shows that the majority of the IR sources within 30" of this candidate are highly reddened. N
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Fig. 4 (Lejt): Adaptively smoothed X-ray image of the DB00-4, 5 , 6 candidate cluster. Circular regions denote the nominal 30" edge of the cluster, and small ellipses show X-ray sources that are possibly correlated with 2MASS sources. (Right): 2MASS K , image of the same region with the same X-ray/IR regions overlaid.
i
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r,,,,, where r,na2is the maximum distance (projected on the plane of the sky) with respect to the barycenter of the cluster for all of the cluster stars. With this process, three possible distribution were generated (hereafter models M I , M2 and M3). In Figure 1, we show the stellar positions for these three modcls.
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3 Numerical simulation Considering the initial conditions given in $2, we have carried out three-dimensional gasdynamic numerical simulations, using the YGUAZU-A binary adaptive grid code. This code integrates the gasdynamic equations with a second order implementation of the “flux-vector splitting” method of van Leer (I982), and is described in detail in Raga, Navarro-Gonzilez & Villagran-Muniz (2000). Due to the fact that the radiative losses are unimportant in this problem (Cant6 et al. 2000), they are not included in the calculations. The numerical simulations were carried out in a five-level binary adaptive grid, with a maximum resolution of 1 . 1 7 2 ~10l6 cm along the three coordinates, covering a cubic domain of 3 x 1018 cm sides, centered on the barycenter of the stellar cluster. This domain was initially filled with an ISM of temperature of lo7 K and density of lo P 3 r r i H cmP3, except for spherical regions of radius R , = 5 x loL6cm, which are centered in each of the cluster star positions. In these spherical regions, spherically symmetric and constant velocity winds are imposed (with temperatures of lo6 K and the mass-loss rate and terminal wind velocity given in $2).
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Raga et al. (2001) computed three models ( M I , M2 and M3). Steady flow configurations are achieved for these three models after a time integration of 1000 yr. In this work, we are going to focus our attention on the evolution of model M1. As an example of the resulting flows, in Figure 2 we show the density and temperature stratification on the y = 0 plane, for model M1 at t = 1000 yr. From this Figure, we see that the central region is filled by an almost homogeneus medium, with densities of l U p 2 ' g cm-3 and temperatures of 5 x lo7 K. This medium was called "the cluster wind' by Cant6 et a1.(2000). Cooler stellar wind cavities are embedded within this cluster wind. Dcnsity and temperature stratification maps (left and right panel of Figure 2, respectively) shows that the central regions have complex morphologies, resulting from multiple stellar wind interactions. However, in some parts o f the periphery, simpler "stellar wind bow shock" structures are observed. These structurcs are the rcsult of the interaction between the cluster wind (generated by the central stars) and the stellar winds of stars located in the periphery.
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Results Predicted X-ray emission
From the thrcc-dimensional numerical simulation of the Archcs cluster wind. we have carried out prcdictions of the X-ray spectrum cmitted in the 0.5-8 keV photon energy rangc. The wavelength-dependent emission coefficient was computed using the CHIANTI' atomic data set (see Dcre et a1.2001 and refercnces therein) as a function of the gas tempcrature. For these calculations, Raga et a1.(2001) assumed that the gas i s in coronal ionization equilibrium and in the low density-regime. They also assumed that the abundance of heavy elemcnts is twice of the solar abundance (following Baganoffet al. 2001). I
The CHIANTI database and associated IDL procedures, are freely available at the following addresses on the World Wide Web: http://www.solar.nrl.navy.mil/chianti.html,http://www.srcetri.astro.it/science/chaint~chiaiiti.html, and http://www.dain~p.cam.ac.ukluser/astro/chianti/cliianti .litml
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With this emission coefficient, Raga et aL(2001) generated X-ray emission maps, integrated over photon energies from 0.5 to 8 keV, in order to directly compare with X-ray images obtained by Yusef-Zadeh et a1.(2002), with the Chandra satellite. For model M1, the resulting map is shown in Figure 3. This simulated map has pointlike intensity peaks, corresponding to the inner stellar wind regions. However, most of the flux comes from a diffuse component, which exhibits several local maxima and filamentary structures. Figure 4 shows the spectrum of the emission integrated over the whole volume of the computational domain. Raga et a1.(2001), applied a frequency-dependentextinction, corresponding to an N H = 1.24 x H cmP3 for directly comparing with observational spectra (Yusef-Zadeh et a!. 2002).
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Finally, from either predicted images o r predicted spectra, Raga et al. (2001) compute a total X-ray 0 ~s-l. ~ luminosity in the 0.5-8 keV band of 2 . 9 ~ 1 erg
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Raga ct a1.(2001) studied the final stationary stage of a cluster wind evolution. However, i t is interesting to analyze the early stages of the evolution. Individual wind sources generate stellar wind bubbles, whosc
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shock waves start to sweep up the circumstellar gas. Interacting regions have high temperatures and pressures, which are good tracers for following the early evolution of these shock waves. In Figure 5, we show maps of thc temporal evolution of the column temperature CT = T dz. At t =20 yr and 60 yr frames, it is seen that the central cluster region is characterized by a filamentary structure, which is made from the superposition of multiple rings. This morphology i s the result of multiple wind interactions, and resembles structures observed in some nebulae, such as the “Honeycomb Nebula” (Meaburn et a1.2001). The filamentary structure evolves to more or less stationary and more homogenous temperature distribution, with some “rays” or “fingers” far away from the cluster harycenter.
1
5
Conclusions
Three-dimensional numerical simulations were carried out for modeling the Arches cluster wind. In order to compare with recent X-ray observations (Yusef-Zadeh et al. 2002), obtained with the Chandra satellite, simulated X-ray emission maps and spectra were generated from the numerical results, and also employing the CHIANTI database (see Dere et al. 2001). Althought the simulated X-ray maps are only qualitatively similar to the observed images (this strongly depends of the employed distribution for the z - coordinate), the X-ray spectra generated for three models (MI, M2, and M3, Raga et al. 2001) are in excellent quantitative agreement with observations. For the erg s-’ was obtained, which is very close to three computed models, a total X-ray luminosity of 3 x erg spl value measured by Yusef-Zadeh et al. (2002) for the Arches cluster. the 4 x We have also analyzed the early cluster wind evolutionary stages, finding that the shock waves from individual stellar wind sources have filamentary structures, which are similar to the structures observed in H a images of the “Honeycomb nebula”. However, the high temperature interaction regions in our simulations do not show appreciable optical line emission. For different model parameters (e.g. a denser initial intra-cluster medium and/or lower velocity winds), the filaments in the initial flow configuration could easily be radiative, and therefore would be observable in optical recombination and collisionally excited lines. Acknowledgements PV and AR are supported by CONACYT grants 34566-E and 36572-E, and DGAPA-UNAM grant INI 12602. LFR acknowledges the support of DGAPA. UNAM, and CONACYT, MBxico. We thank Israel Diaz for computer support.
References Baganoff, B. K., Bautz, M. W., Brdndt W. N. et al. 2001, Nature, 413, 45 Cantb, J., Raga, A. C. & Rodriguez, L. F. 2000, ApJ, 536,896 Cotera, A. S., Erickson, E. F., Colgan, S. W. J., Simpson, J. P., Allen, D. A., & Burton, M. G. 1996, ApJ, 469, 729 Dere, K. P., Landi, E., Young, P. R., & Zanna, G. 2001, ApJS, 134,331 Figer, D. F., Kim, S. S., Moms, M., Serabyn, E., Rich, R. M., & McLean, I. S. 1999, ApJ, 525, 750 Lang, C. C., Goss, W. M. &Rodriguez, L. F. 2001, ApJ, 551, L143 Meaburn, J., Redman, M. P., Bryce, M., Lbpez, J. A., Al-Mostafa, 2. A., & Dyson J. E. 2001, ApSS, 272, 217 Nagata, T., Woodward, C. E., Shure, M. & Kohayashi, N. 1995, AJ, 109, 1676 Osernoy, L. M., Gcnzel, R. & Usov, V. 1997, MNRAS, 288,237 Raga, A. C., Velazquez, P. F., Cantb, J., Mdsciadri, E. & Rodn’guez, L. F. 2001, ApJ, 559, L33 Serabyn, E., Shupe, D., & Figer, D. F. 1998, Nature, 394, 448 van Leer, B. 1982, ICASE Rep., 82-30 Yusef-Zadeh, F., Law, C., Wardle, M., Wang, Q. D., Fruscione, A., Lang, C. C., Cotera, A. S. 2002, ApJ, 570, 665
SiO Maser Sources within 30 pc of the Galactic Center
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Shuji Deguchi* and Hiroshi Imai ' Nobeyama Radio Observatory,Minamimdki, Minamisaku, Nagano 384-1305. Japan VERA project office, National Astronomical Observatory, Mitaka, Tokyo 18 1-8588, Japan present address: JIVE, PO Box 2, 7990 AA, Dwingeloo, the Netherlands
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Key words Masers, Galactic Center, AGB, Mass-loss Abstract. Using the Nobeyama 45-111radio telescope, we have observed 314 large amplitude variables within 30 pc of the Galactic center in SiO maser lines. Resulting detections give the radial velocities of 174 stars; light-variation periods have been known for all of thcse stars. The SiO detection rate increases sharply with the period and it is about twice of the OH maser dctection rate. The radial-velocity data show slow and rapid rotations of the outer and inner circumnuclear-diskstars, respectively. Five high-velocity stars were found only at the negative-longitudeside of the Galactic center. Estimation of the ages of high velocity stars suggests that these stars must be be accelerated to high velocities within 10' years.
1 Introduction Mass-losing AGB stars are good probes of the Galactic center. For the last few years, we have observed steller mascr sources near the Galactic ccnter in SiO maser lines at 43 GHz (Deguchi et al. 2000) with the Nobeyama 45-m radio telescope. These observations were made toward color-selected IRAS sources. Because of the incompleteness of the IRAS survey near the Galactic center, the area within 1 degree from the Galactic center had remained unsurveyed i n SiO masers. The situation, however, has recently changed owing to new ground-based surveys with near-infrared array cameras, and to space-based mid-infrared surveys. Large numbers of candidate stars suitable for maser surveys toward the nuclear disk have been discovered in the near-infrared K band by Glass et al. (2001) making use of their characteristic largeamplitude variability. In the present paper, we report the SiO maser survey of the large amplittide variables within 30 pc of the Galactic center. Resulting detections gave the radial velocities of 174 stars. These sources are intermediate-mass stars in the Asymptotic Giant Branch phase of the ages between lo7 - 10"' y. Based on these data, we discuss on kinematics of masers stars in the Galactic center.
2 Observations and results Observations in the SiO .J = 1 - 0 o = 1 and 2 transitions at 42.821 and 43.122 GHz were made with the 45-m radio telescope during February 2001 - May 2002 i n a long-term project of the Nobeyama radio observatory. A detail description of the 45-in telescope system can be found in Izumiura et al. (1999). A preliminary result of the survey of 134 large amplitude variables in the Galactic center was published i n Imai et al. (2002). Monitoring observations of SiO sources towards Sgr A* were reported in Deguchi et al. (2002). Examples of SiO maser detections are shown in Figure 1. For most of the SiO maser sources, the intensity ratio of the SiO J = 1 - 0 'r = 1 to '1' = 2 line is near unity. It is know to weakly correlate * Corresponding author: e-mail: deguchiQnro.nao.ac.jp, Phone: +81 267 984369, Fax: +81 267 98 2884
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with the IRAS 12/25 pm color (see discussion in Nakashima and Deguchi 2003). Therefore, simultaneous observations of two SiO lines secure detections. The 45-m telescope beam size is about 40” (HPFW). Because of the high source density at the Galactic center, we occasionally had double or triple detections of sources in the same beam (as shown on the right in Figure 1). Even in such cases, we can assign the source to a particular object by observing an offset position from the source of 10-15” and measuring the relative intensity variation according to the position. We have observed 314 objects, resulting in 174 detections through June 2002.
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Fig. 1 Spectra of SiO masers toward g6-247 (left) and g3-46 (right). The right shows a case of triple detections at KSr = -141, -98, and -45 km s-’. The number under the source name indicates the observed date in yymmdd.d
format.
2.1 Period-Detection rate Figure 2 shows a histogram of the detection rate versus variablity period. The SiO detection rate is approximately 55 %, which is similar to the SiO detection rate for the inner bulge IRAS sources. This fact suggests that the sample of the large amplitude variables in the Galactic center has very similar characteristics to the color-selected bulge IRAS sample. The SiO detection rate sharply increases with the period of light variation. We note that the average period of the Glass et al. (2000) sample is about 430 days. Because the present sample is considered to be at a uniform distance, a relation between the detection rate and the period is firmly established. Blind OH maser surveys have been made with VLA (Sjouwerman et al. 1998) in the same region of the sky. The SiO detection rate i s about twice of the OH detection rate. 2.2 Longitude-Velocity Diagram Figure 3 shows a longitude-velocity (I-u) diagram. The best fit to the velocity data is = 0.0594(Al/‘/) - 6.85 km sC1
(1)
for all of the sources, and
V,,,
= 0.274(Al/’/)
- 2.90
km s-l
(2)
for sources within 300”. The radial velocity increases with the longitude offset, Al, from Sgr A * due to Galactic rotation; on average, the rotational rate is about 214 km lip’ per degrce, which is close to the Galactic rotation obtained from CO and HI observations (Honma and Sofue 1997). However, the rotational velocity tends to increase near the Sgr A* (within 5’) as shown in equation (2) (see also Deguchi et al. (2002) within 2’). The overall velocity structure of the SiO 1-ZJ diagram is similar to the OH b-IJ diagram (for example, see Sjouwerman et al. 1998). We can recognize a hole in the SiO I-,ii diagram at
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Period(d) Fig. 2 histogram of the period. Shaded area shows SiO detection and blank nondetection. The line graph shows the detection rate (the scale on the right axis), and the broken lines shows the OH detection rate for the same sources in the present sample.
= 200", = 50 km s-'), which i s also seen in the OH l-r~ diagram. It i s possible that this hole is produced by some non-axisymmetric gravitational potential in the Galactic center.
(Al
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Fig. 3 Longitude-velocity diagram for stars with periods below and above 600 d. The leasf-square tits to each subsample were shown.
S. Deguchi and H. Irnai: SiO Masers within 30 pc of GC
288
Figure 3 also exhibits a slight asymmetry in velocity dispersion with respect to Sgr A*; on the A1 < 0 side, the points seem scattered more than those on the A1 > 0 side. To check the asymmetric structure in velocity dispersion, we computed the dispersions from the best fit lines for the subsets inside and outside of r = 300”. The results were summarized in Table 1. The high velocity stars (JV,, 1 > 200 km s-l) strongly influence to the results. Therefore, we calculated the cases including and excluding the high-velocity stars. Even excluding the high velocity stars, the velocity dispersion is significantly larger on the Al < 0 side. Table 1 Velocity Dispersions.
A1 > 0, r > 300‘‘
A1 < 0, r > 300”
r < 300’’
(km s-l)
(km s - * )
(km s-l)
All sources
49.4
90.1
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Excluding high-vel. sources
49.4
61.2
57.9
Sample
To check the presence of period (or mass) segregations of stars in the stellar cluster, we also computed the velocity dispersions for subsets o f stars with periods below and above 600 days. No period segregation was found in the present sample, indicating that the mixing time scale in phase space is sufficiently larger than the ages of these stars. It is notable that periods of stars circulating around the Galactic center are smaller than 4 x 106 years within 30 pc of the Galactic center. Even for high progenitor-mass stars (9 A f @ ) , the ages in the AGB phase (for example, Vassiliadis & Wood 1993) are longer than this dynamicalmixing time scale. 2.3 High-velocity Sources Figure 4 shows SiO maser spectra and 2MASS J H K images of 5 high velocity sources. These high velocity sources appear at both positive and negative velocities, therefore the previously suggested asymmetry in velocity (van Langevelde et al. 1992), can be regarded as a statistical fluctuation. It is possible that Lhese high velocity sources have highly elongated eccentric orbits around the Galactic center. From the luminosities and periods, we can estimate ages of these stars. The mass of the brightest star, g3-2855, is estimated to be more than 5 A&. The age must be shorter than lo8 y. Therefore, this high velocity star must be accelerated to V 300 km/s within lo8 y. Also, the velocity distribution of high velocity stars seems to be discreet from the low velocity stars (for the inner bulge stars ((!I > lo), the distribution of high velocity stars seems continuous to the distribution of low velocity stars). The time scale of stars circulating around the Galactic center is about 4 x lo6 years. The ages of these stars are longer than 2 x lo7 years (for M < Shf,), so the distribution of these high-velocity stars should be smoothed out enough as a result of rapid phase mixing. Therefore, the fact that these high velocity stars were found only at A1 < 0 side is unexplained. Kim and Morris (2001) proposed that vertical diffusion of nuclear disk stars is driven by the scattering of stars o f fgiant molecular clouds in the nuclear disk. This time scale is longer than lo9 y, indicating that this mechanism cannot he used to explain observed high velocity stars. An efficient mechanism to disrupt thc star cluster falling into the central blackhole was proposed in this conference by Kim, Morris, & Figer, which may possibly create the high velocity stars near the Galactic center in a free fall time scale.
-
2.4
Around Sgr A*
Figure 5 shows a time variation of SiO intensities of the sources which are located within 20” from Sgr A*. By mapping the 100” x 200” area of Sgr A*, we detected about 15 SiO sources (Deguchi et al. 2002), and obtain the positions in 5-10’’ accuracies. We detected one high velocity SiO source (-342 km s p l component; No. 5 in the lower panel of Figurc 5). This component was more clearly detected in 1997 N
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-350
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'
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Fig. 4 SiO maser spectra and false-color J H A ' images of.5high velocity sources. The high velocity source found toward G4-1 13 is different from G4-I 13 (OH 359.855-0.078) with Vb,. = 5.2 km s - ' ; in order to clarify the difference, additional number (.2)was attached.
(Izumiura et al. 199X). Because of it's intensity, it was quite difficult to locate the position of this star i n the mapping observations, but is estimated to be no further than 10" away from Sgr A*. The -27 km s pl component, which was attributed to IRS IOEE located about 10" NE of Sgr A" (Menten et al. 1997), showcd a large time variation. The SiO masers of IRS lOEE flarcd in March - May 2000 and again in March 2002. We have monitored the SiO maser intensity of this object for last few years. The result of this monitoring till 2001 was reported in Deguchi et al. (2002). The intrinsic maser intensity (about 1.5 Jy) of IRS lOEE SiO masers during the flares is approximately thc same as those of
S. Deguchi and H. Imai: SiO Masers within 30 pc of GC
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the strongest SiO maser sources in our Galaxy: Orion SiO masers and Sgr B2 MD5. This fact indicates that the maser power at the flare time is comparable to the maximum seen i n SiO masers.
0.4 0.3
-s
0.2
(d
-I
0.1
0 '
-0.1 ~ ~ " " " " " " ' ~ " " " " " " " " " ' ~ -400 -300 -200 -100 0
V
Isr
100
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'
" ~ 400
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2
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Fig. 5 Time variation of the SiO maser sources around Sgr A* (within 20"). Upper panel is the spectra taken on June 9, 1999 and the lower panel on May 25, 2000. The SiO masers from IRClOEE at -27 km s C 1 flared during March-May 2000.
Near-infrared time variation of IRS lOEE was monitored by Wood et al. (1998) and the light-variation period of this star is 720 days. It was not certain whether the SiO maser flares in this period or not. SiO masers of this star flared again in March 2003. Therefore, they seem to flare in every two years. N
Astron. NachrJAN 324. No. SI (2003)
29 1
3 Conclusions We found that the SiO detection rate increases sharply with the period and it is about twice of the OH maser detection rate. SiO masers are useful tool to investigate the relatively low-mass AGB stars in the Galactic center region. The radial-velocity data show slow and rapid rotations of the outer and inner circumnuclear-disk stars, respectively. No period/mass segregation was found in the sample, indicating that a phase mixing occurs in a dynamical time scale of about lo6 y. Five high-velocity stars were found only at the negative-longitude side of the Galactic center. SiO masers of IRS IOEE were found flared i n May 2000. This SiO maser flare seems to occur periodically in every two years. Acknowledgements The observations were made for a long-term project of Nobeyama Radio Observatory with collaboration of Drs. T. Fujii, , I . S. Glass, Y. Ita, H. Izumiura, 0. Kameya, A. Miyazaki, Y. Nakada, and J. Nakashima. This research was partly supported by Scientific Research Grant (C2) 12640243 of Japan Society for Promotion of Sciences.
References Deguchi, S., Fujii, T.,Izumiura. H., Kameya, O., Nakada, Y., Nakashiina. J., Ootsubo, T., & Ukita, N., 2000, ApJS, 128,571 Deguchi, S., Fujii, T., Miyoshi, M., & Nakashima, J. 2002, PASJ, 54, 61 Glass, I. S., Matsumoto, S., Carter, B. S .. & Sekiguchi, K. 2001, MN. 321,77 Honma, M., & Sofue, Y. 1997. PASJ, 49,539 Imai, H., Deguchi. S. Fujii, T., Glass, I. S., Ita, Y., Izumiura, H., Kaineya, O., Miyazaki. A,, Nakada, Y., & Nakashima, J.. PASJ, 54, 19 Izumiura, H., Deguchi, S., & Fujii, T. 1998. ApJ. 494, L89 Izumiura, H.. Deguchi, S., Fujii, T., Kameya, 0..Matsumoto, S., Nakada, Y., Ootsubo, T., & Ukita, N. 1999, ApJS, 125,257 Kim, S.S. & Moms, M. 2001, ApJ, 554, 1059 Menten, K. M., Reid. M. J., Eckart, A., & Genzel, R. 1997, ApJ. 475, LI I 1 Nakashima, J., & Deguchi, S. 2003 PASJ, 55, No. I , in press Sjouwerman, L. 0..van Langeveldc, H. J., Winnberg, A,, & Habing, H. J. 1998, A&AS, 128,35 van Langevelde, H. J., Brown, A . G. A.. LindqvisL M., Habing, H. J., de Zeeuw, P. T. 1992, A&A, 261, L17 Vassiliadis, E., & Wood, P. R., 1993, ApJ, 413,641 Wood, P. R., Habing, H. J.. & McGregor, P. J. 1998, A&A, 336, 925
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Astron. Nachr./AN 324, No. S 1,293-297 (2003) / DO1 10.1002/asna.200385037
86 GHz SiO masing late-type stars in the Inner Galaxy M. Messineo* I, H.J. Habing', L.O. Sjouwerman', K.M. Menten', and A. Omont4
' Leiden Observatory, P.O. Box 9513,2300 RA Leiden, the Netherlands
' National Radio Astronomy Observatory, P.O. Box 0, Socorro NM 87801, USA Max-Planck-Institut fur Radioastronomie,Auf dem Hugel 69, D-53121 Bonn, Germany Institut d'Astrophysique de Paris, 98bis Boulevard Arago, F 75014 Pans, France
Key words Masers, Circumstellar matter, Galaxy: kinematics and dynamics
We present 86 GHz ( v = 1,J = 2 t 1) SiO maser line observations with the IRAM 30-ni telescope of a sample of late-type stars in the inner Galaxy (30" < 2 < -30'). The stars were selected from the TSOGAL and MSX catalogues on the basis of their mid-infrared fluxes and colours. SiO maser emission was detected towards 268 (6 1%) of our targets, thereby doubling the number of maser line-of-sight velocities measured toward the inner Galaxy. Our sample consists mostly of Mira-like stars. They are more numerous than OWIR stars which were previously observed to measure line-of-sight velocities. The revised longitudevelocity diagram of the inner Galaxy clearly shows a stellar nuclear disk.
1 Introduction Asymptotic Giant Branch (AGB) stars are good tracers of the Galactic structure and kinematics. Being bright in the infrared, they are visible in directions with high extinction. Maser emission from their circumstellar envelopes is strong enough to be detected throughout the Galaxy and reveals the line-of-sight velocity of the stars to within a few km s-' . Frequently detected maser lines are from O H at 1.6 GHz, HzO at 22 GHz, and SiO at 43 GHz and 86 GHz (e.g. Habing 1996). Until recently only a few hundred stellar line-of-sight velocities were known towards the inner regions of the Milky Way ( 30 O < 1 < -30 O and (bl < 1).These are mainly of OWIR stars, which are AGB stars with OH maser emission in the 1612 MHz line, mostly undetected at visual wavelengths. This number is too small to allow for a good quantitative multi-component analysis of Galactic structure and dynamics. Therefore, obtaining more line-of-sight velocities remains an issue of prime importance. To enlarge the number of known stellar line-of-sight velocities, using the IRAM 30-m telescope we have started a survey for 86 GHz S i O (v = 1,J = 2 + 1) masers towards an infrared-selected sample of late-type stars.
2 Source selection A large variety of names exists to indicate oxygen-rich AGB stars characterized by different pulsation properties and/or mass-loss rates: semi-regular (SR) stars and Mira stars (visual pulsation amplitude larger than 2.5 mag), large amplitude variables (LAV), long pcriod variable (LPV) stars (when their periods are longer than 100 days), and OH/IR stars (with 1612 MHz OH maser emission). In the IRAS colour-colour diagram the oxygen-rich AGB stars are distributed on a well-defined sequence of increasing shell opacity and stellar mass-loss rate (e.g. Habing 1996) which goes from Miras with the bluest colours and the 9.7 p m silicate feature in emission, to O M R stars with the reddest colours and the 9.7 pm silicate feature in absorption. * Corresponding author: e-mail: messineoQstrw.leidenuniv.nI, Phone: +31 71 5275831, Fax: +31 71 5'275819
02003 WILLY-VCH Vrrl.ig GmbH & Co
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Fig. 1 Location of the observed stars, irrespective of detection and non-detections, in Galactic coordinates. The MSX sources are shown as open squares and the ISOGAL sources as filled circles.
SiO maser emission is generated in the envelopes of mass-loosing AGB stars, close to their stellar photospheres. It occurs more frequently towards oxygen-rich Mira stars than towards other AGB stars. Furthermore, the relative strengths of different SiO maser lines are observed to vary with AGB type (Nyman et al. 1993 and references therein). It indicates that the SiO maser properties depend on the stellar mass loss rate and on the stellar variability and that, differently than for OWIR stars, for Mira-like stars the 86 GHz (v = 1, J = 2 + 1) SiO maser transition is a good tool to obtain stellar line-of-sight velocities. Therefore, we selected Mira-like stars. They also are far more numerous than OHlIR stars. The stars to be searched for maser emission were selected from both a preliminary version of the combined ISOGAL-DENIS catalogue (Schuller et al. 2003; Omont et a]. 2003) and from the MSX catalogue (Egan et at. 1999). ISOGAL is a 7 and 15 pm survey of -16 deg2 towards selected fields along the Galactic plane, mostly toward the Galactic centre. The Midcourse Space Experiment (MSX) is a lower sensitivity and resolution survey covering the entire Galactic plane at five mid-IR bands ranging from 4.3 pm [Bl band], to 21.4 pm [ E band]. The SiO maser search was limited to the Galactic plane between (b( < 1 O and I = +30° and 2 = -4"; the lower limit in longitude is imposed by the northern latitude of the IRAM 30-m telescope. The brightest 15 pm sources, with a magnitude [I510 < 1.0, and those with a 7ym-15pm colour, ([7]0 - [15]0) < 0.7, were excluded since they are likely to be foreground stars. Because of the general correlation of SiO maser emission and IR luminosity (Bujarrabal et at. 1987), sources with [15]0 > 3.4 were excluded since they are likely to show SiO maser emission fainter than our detection limit of 0.2 Jy. Sources with ( [ 7 ] 0- [15IO) > 2.3 were excluded since they are likely to be compact HI1 regions or young stellar objects. Furthermore, for the ISOGAL sample a range of intrinsic (Ks0- [15]0) colour was selected to avoid thick envelope OHAR stars and young stars. No similar restriction could be applied to the MSX sample because no near-infrared counterparts were available at the time of the observations. As an alternative to the ( K s o- [15]0) criterion, for the MSX sample we imposed an upper limit to the ratio of the fluxes in band E and C (FE/Fc< 1.4). For the MSX sources additional variability information was taken into account. Moreover, sources close to a known OH maser were discarded as the kinematic data are already known.
3 Results The observations were carried out with the IRAM 30-m telescope (Pico Veleta, Spain) between August 2000 and September 2001. With a detection limit of 0.2 Jy we detected SiO (,u = 1,J = 2 + 1)maser emission towards 268 stars, of which 255 were previously unobserved. In almost all cases the targeted star
Astron. Nachr./AN 324, No. S1 (2003)
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is the only mid-infrared object falling inside the main beam (29"). The total detection rate is 61%. The SiO line width distribution ranges between 2 and 16 km s-l with a peak at 4 km s-'. We observed 15 LPVs found by Glass et al. (2001) and detected SiO maser emission from 1 1 of them (73 %). Since the observations were taken at a random pulsation phase, and since the SiO maser intensity is known to vary during the stellar phase by up to a factor ten (Bujarrabal 1994), this detection rate is a lower limit to the actual percentage of LPV sources characterized by 86 GHz SiO maser emission. Only 16 % of those LPV stars have associated OH emission (Glass et a]. 2001). and only 23 % among those within our defined colour-magnitude region. Among large amplitude variable AGB stars, the 86 GHz SiO masers are much more frequent than OH masers.
4 Infrared measurements The combination of near- and mid-infrared photometry permits us to study the nature of the stars, to derive luminosities, mass-loss rates, and provides a good discrimination of foreground stars. We searched for possible counterparts of our SiO targets in the recent extensive infrared data catalogues: DENIS, 2MASS, ISOGAL and MSX. We found (Messineo et al. in preparation) that our stars are intrinsically very bright in the near-infrared, and not heavily obscured by circumstellar matter, unlike the OWIR stars which are often obscured even in the K s band (e.g., Ortiz et al. 2002). Using near and mid-infraredphotometry it is possible to calculate the interstellar extinction and apparent bolometric magnitudes. This will add information on the source distances and improve the understanding of the Galactic longitude-velocity diagram. Therefore, 86 GHz SiO masers from Mira-like stars are better tracers of the Galactic structure than OWIR stars.
2 -
18
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12 10 J (DENIS)
8
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0 -
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Fig. 2 Left Panel: The difference between the 2MASS and DENIS J magnitudes versus the DENIS J magnitude. Squares show the position of our ISOGAL SiO targets. For comparison small dots show ISOGALlDENlS and 2MASS associations obtained in several ISOGAL fields. No correction for offset in the photometric zeropoint has been applied. Right Panel: The difference between the 2MASS and DENIS J magnitudes versus the difference of the 2MASS and DENIS K s magnitudes. The crosses indicate 2MASS J upper limits.
4.1
Variability
The 86 GHz SiO maser line intensity is observed to be stronger in oxygen-rich Mira stars. To increase the chance of detecting SiO maser emission, we selected MSX sources with evidence for variability. Unfortunately, the ISOGAL database does not provide information on variability. However, from their position
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in the ISOGAL/DENIS (Kso - [15]o) versus [15]0 colour-magnitude diagram, it was expected that the ISOGAL SiO targets would be mainly large amplitude variables (Messineo et al. 2002). The J and K s filters used by DENIS and 2MASS are similar, therefore the measurements obtained during the course of the DENIS and 2MASS surveys are directly comparable. For each of the 61 ISOGAL fields containing SiO targets, we cross-correlated the ISOGAL/DENIS and the 2MASS point source positions. For 55% of our ISOGAL stars, the difference between the 2MASS and the DENIS J and the 2MASS and the DENIS K s is larger than 3 times the field dispersion (see Fig. 2). Therefore, our sample contains mostly variable stars. Due to the simultaneity of the J and K s measurements in both the DENIS and 2MASS surveys, a correlation is expected and found between the variation in the J magnitude (4J)and in the K s magnitude ( A K s ) shown , in Fig. 2. A linear least squares fit gives,
AJ
= 1.71(&0.03) x AlCs
+ 0.02 i (0.02).
-
60% of (he amplitude in J band. This results is The relative pulsation amplitude in the K s band is in agreement with the relation between the pulsation amplitudes in J and in K s bands for oxygen-rich Mira stars in the solar neighborhood (Olivier et a]. 2001). A monitoring program of our SiO maser sample will provide pulsation periods and will yield also an estimate of the Galactic center distance through the period-luminosity relation.
5
Longitude-velocity diagram
The line-of-sight velocities of the 86 GHz SiO maser sources range from -274 to 300 km spl, which is consistent with previous stellar maser measurements and with the l2COvelocities toward the inner Galaxy (Fig. 3). An appreciable number of SiO sources are located in a region forbidden for pure circular rotation, at negative velocities between O"< I < 20". Around zero longitude the stellar distribution follows the high velocity gas component of the nuclear disk. Nuclear disk stars are heavily extinct ( A v > 20 mag, Messineo et al. 2003), as seen from Fig. 4, and their line-of-sight velocities range from 150 to -200 similarly to the gaseous nuclear disk line-of-sight velocities (cf. the "CO (1 - ,u) diagram in Fig. km SKI, 4 of Bally et al. 1988). The stellar nuclear disk (Ill < 1.5; Ibl < 0.5; A v >20 mag) rotates very rapidly around the Galactic center: our best-fit slope is 175(+40) km s-l per degree, consistent with the value (180 f 15 km s-l) found for the OHRR stars in the Galactic center (Lindqvist et al. 1992).
-
References Bally, J., Stark, A. A,, Wilson, R. W., and Henkel, C. 1988, ApJ, 324, 223 Bujarrabal, V., Planesas, P., and del Rornero, A. 1987, A&A, 175, 164 Bujarrabal, V. 1994, A&A, 285,953 Dame, T. M., Hartmann, D., and Thaddeus. P. 2001, ApJ, 547, 792 Egan, M. P., Price, S. D., Moshir, M. M., et al. 1999, AFRL-VS-TR-1999, 1522 Fux, R. 1999, A&A, 345,787 Glass, I. S., Matsurnoto, S., Carter, B. S., and Sekiguchi, K. 2001, MNRAS, 321, 77 Habing, H. J. 1996, A&A Rev., 7, 97 Lindqvist, M., Habing, H. J., and Winnberg, A. 1992, A&A, 259, 1 I8 Messineo, M., Habing, H. J., Sjouwerrnan, L. O., Omont, A,, and Menten, K. M. 2002, A&A, 393, 115 Messineo, M., Habing, H. J., Ornont, A., Menten, K. M., and Sjouwennan, L. 0.2003, in preparation Nymdn, L.-A,, Hall, P. I., and Le Bertre, T. 1993, A&A, 280, 551 Olivier, E.-A,,Whitelock, P., and Marang, F. 2001, MNRAS, 326, 490 Omont, A,, Gilrnore, G. F., Alard, C., et al. 2003, A&A, 403,975 Ortiz. R., Blommaert, J. A. D. L., Copet, E., et al. 2002, A&A, 388,279 Schuller, F., Ganesh, S., Messineo, M.. et al. 2003, A&A, 403, 955
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Fig. 3 Stellar longitude-velocity diagram overlayed on the grayscale CO (1 - v) diagram from Dame et al. (2001). The SiO 86 GHz masers are shown as dots. Gas features are labelled following Fig. I of Fux (1999).
-2
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Galactic longitude [degrees]
Fig. 4 Stellar longitude-velocity diagram of our 86 GHz SiO masers. The filled circles indicate sources with visual interstellar extinction larger than 20 mag. Most of those clearly helong to the nuclear disk, fast rolating component. The continuum line indicates our best fit to the nuclear disk component.
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Astron. NachrIAN 324. No. S1. 299-302 (2003) / DO1 10.1002/asna.200385083
CNO Abundances in the Quintuplet Cluster M Supergiant 5-7 S. V. Rarnirez*’, K. Sellgren2, R. Blum’, and D. M. Terndrup’ ’ Infrared Processing and Analysis Center, CaliforniaInstitute of Technology, Mail Code 10@22,770 South Wilson Avenue, Pasadena, CA 9 I 125, USA Department of Astronomy,The Ohio State University, 140 West 18th Avenue, Columbus, OH 43210 USA NOAO/CTIO, Casilla 603, La Serena, Chile
Key words stars, abundances, Galactic center
Abstract. We present and analyze infrared spectra of the supergiant VR 5-7, in the Quintuplet cluster 30 pc from the Galactic center. Within the uncertainties, the [CIH], “/HI, and [OfH]abundances in this star are equal to those of a Ori, a star which exhibits mixing of CNO processed elements, but are distinct from the abundance patterns in IRS 7.
1 Introduction We have previously published a differential analysis of the iron abundance [Fe/H] in ten cool, luminous stars within 30 pc of the Galactic Center, compared to 1 I stars of similar temperature and luminosity in the solar neighborhood. We found that both samples of stars had a narrow distribution of [Fe/H] centered at the solar value (Ramirez et al. 2000). Carr et al. (2000) also studied [C/H], “/HI, and [O/H] in (1 Ori and in IRS 7, which lies at a prqjected distance of 0.25 pc from Sgr A*. They found that dredge-up of CNO-processed material was present in both M supergiants, as expected theoretically, but that the amount of internal mixing was much stronger in IRS 7, stronger than predicted by current evolutionary models. Our main goal is to investigate whether the strong internal mixing found in IRS 7 is due to the unusual conditions for star formation in the central 100 pc of the Galaxy (Morris 1993, Morris & Serabyn 1996),or is due to some tidal interaction between IRS 7 and the supermassive black hole (2.6 x lo6 Ma: Schoedel et al. 2002, Ghez et al. 2000) at the Galactic Center. Here we present preliminary results on [ C k l ] ,[ N/H], and [O/H] in the M supergiant VR 5-7, which lies in the Quintuplet cluster, 30 pc from the Galactic Center.
2 Observations and Analysis High-resolution (A/AA 25,000) K-band and H-band spectra of VR 5-7 were obtained through Gemini sponsored access to the Keck Telescope using NIRSPEC in June 2001. The spectra were reduced using IRAF, involving Rat fielding, sky subtraction, spectrum extraction, wavelength calibration, and removal of atmospheric absorption features. The observed spectra are of high quality, with a signal to noise ratio above 100, which is needed for detailed abundance analysis. The abundance analysis was done using a current version of the LTE spectral synthesis program MOOG (Sneden 1973). The program requires a line list with atomic and molecular parameters and an input model atmosphere for the effective temperature and surface gravity appropriate for the star. The atomic and molecular parameters (wavelength, excitation potential, gf-value, damping constant, and dissociation constant) were obtained the same way as in Ramirez et al. (2000), and also included C O molecular parameters N
* Corresponding author: e-mail:
[email protected],Phone: + I 626395 1919, F d X : +I 626397 7018
@ 2002 WILEY-VCH Veiiag GmhH & Co. KCaA. Weiiihrini
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from Goorvitch (1994) and OH molecular parameters from Black (private communication). The solar abundance model atmospheres from Plez (1 992) were used for our abundance analysis. The stellar parameters of VR 5-7 were taken from Ramirez et al. (2000): effective temperature T.fi = 3500 K, surface gravity logg = -0.2, microturbulent velocity = 2.9 km s - ' , and macroturbulent velocity = 12.6 km
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Fig. 3 a) Pact map of the central 2 pc, with the inner 19" comprised of the high resolution Camera 2 data. Diamonds indicate the locations of the broad-line stars while the narrow-line stars are indicated by circles. All of these stars show emission in this Paa filter. b) All emission line star candidates, in the central region, marked by square boxes. The intensity scale is stretched to show fainter emission.
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Astron. NachrJAN324, No. SI, 309-313 (2003) / DO1 10.1002/asna.200385085
Ten Thousand Stars Toward the Galactic Center * Frangois Rigaut**’,Robert Blum2,Tim Davidge3,and Angela Cotera4 ’ Gemini Observatory, 670 N A’Ohoku place, Hilo, HI-96720, USA Cerro Tololo Interamerican Observatory, La Serena, Chile Hertzberg Institude for Astrophysics, Victoria, Canada Steward Observatory, Tucson, USA
Key words Astronomical Instrumentation, Stars, Star formation
Abstract. We report on the Galactic center data set obtained at the Gemini Observatory as part of the demonstration science program with the Hokupa’a adaptive optics system. The data set is presented and characterized. Preliminary results on the stellar populations from CO band data and IRS8 are presented.
1 Introduction To check the performance of the telescope and the Hokupa’a+QUIRC instrument, to test observing procedures and data reduction methodologies and train the science staff, the Gemini Board approved in early 2000 a program of “Demonstration Science” for this instrument. After exchanges with the science community through their representatives at the GSC, the Galactic Center was chosen as the most relevant program. A science group was put together, involving members from the Gemini partner countries. Taking advantage of the strength of Hokupa’a (faint guide star, large imaged field) and keeping its weaknesses in mind (lower compensation capabilities than e.g. the Keck system due to the limited number of actuators), the science group converged toward the following science objectives: (1) Stellar Population: Determine the properties of the stars in function of their position in this region, particularly of their distance to Sgr A*. The diagnostics selected were H, K’, CO and CO continuum images. Because of the limitations of the A 0 system at short wavelengths, J was not selected, (2) Variability: Monitor a couple of fields (Central Sgr A* and Arches cluster) for variable stars, (3) Map the extremely structured extinction, and (4) Creatiodenlargement of an astrometric database over a larger field than previous data sets. 2 GB of data were obtained, covering a total of 4000 square arcsec in the vicinity of Sgr A*. These images have been reduced and were released to the international community in November 2000. The data set is briefly described in this paper, together with preliminary analysis of some remarkable features of this region unveiled by these observations (Rigaut et al 2003). A number of papers have been published to date that use this data set (see references).
2 Instrumentation and Observations The University of Hawaii Hokupa’a A 0 system, developed and operated by the UH A 0 group with the support from the NSF, uses curvature sensing and is able to use rather faint guide stars (nominally down * Based on ObSeNatiOnS obtained at the Gemini Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under a cooperative agreement with the NSF on behalf of the Gemini partnership: the National Science Foundation (United States), the Particle Physics and Astronomy Research Council (United Kingdom), the National Research Council (Canada), CONICYT (Chile), the Australian Research Council (Australia), CNPq (Brazil) and CONICET (Argentina) ** Corresponding author: e-mail:
[email protected],Phone: + I 808 9743686, Fax: +1 808 935 9650
@ 2W3 WILEY-VCH Veilag GmbH d To KGaA, Weinheim
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originally for use on a 4-m class telescope (36 subapertures), and therefore is undersized for use on Gemini in the sense that it has too few actuators for adequate compensation at short IR wavelengths under typical seeing conditions. Reasonable performance in terms of compensated image quality (IQ) can be obtained only under periods of medium to good seeing. The data cover 11 fields of 20” x 20” each. Fields 1 4,12, and 13 are the central field data set. Fields 9 and 10 are control fields to check the properties of the bulge background population. Finally, field I 1 is the Arches Cluster. Fig.1 shows the respective location of the fields. H and K’ images are available for all fields. Dots next to the field number indicate fields for which CO (2.29 pm) and CO continuum (2.26 pm) have been taken. The guide star location for each field is marked with an (orange) star symbol. Guide star magnitude range from 13.5 to 15.1. The data reduction was done in the standard way, by several members of the science team, and images were cross checked to insure the quality and reproducibility of the data reduction process. Depending on the seeing during acquisition, the IQ on the image of this data set varies between 0.085” and 0.18” close to the guide star (median=O.l225”) and 0.11” to 0.23” at the edge of the field (median 0.155”).Crowding in the central part of the cluster goes up to 10 detected stars per square arcsec. Limiting magnitude in this data set varies from K’ x 17 in the most crowded area to K‘ > 21 in the least crowded area.
3 Identification of the young bright stars in the GC core from CO data We report here results from an analysis of the CO data in field 1 and 2 (the field containing Sgr A* and the one immediately to the north of it), followed by a preliminary interpretation of the data. The CO feature is usually a tracer of late-type giant and supergiant stars. By doing these observations, we were trying to discriminate between various stellar population ages. A combination of a CO filter (2.29pm, bandpass=20A) and a K continuum (2.26pm, bandpass=60A), also called CO continuum here, was used. A longer wavelength continuum was not available for these observations. We note that the use of a single continuum filter might affect the estimate of the CO absorption in the case of objects with large color gradients. However, our data indicate that this effect is small compared to the photometric errors. A 0 observations are particularly suited to narrow band photometry, as the observations in the filter and in the immediately adjacent continuum can be taken close enough to each other in time (possibly even interlaced) that the same turbulent conditions apply. The isoplanatic PSF degradation off the guide star is also very similar, considering that the wavelengths are very close. Therefore, even if the error on the absolute magnitude versus the position in the field in the CO filter remains large, for example, the same error is encountered in the continuum filter, thereby compensating for the introduced error. Figure 3 presents the (CO - K continuum = CO index) vs (K continuum) CMD. The CMD for field 2 is shown on the right hand panel. The scatter about a linear fit can be considered as an upper limit estimation of the noise. The scatter on the CO index is 2.5% (0.025 magnitude, 1 sigma) for Kc < 16. To our knowledge, it is one of the first quantitative estimates of the photometric noise in A 0 and, we believe, the lowest. The left hand panel of figure 3 presents the same CO index vs CO continuum, but including field 2 and field 1 -the field containing Sgr A*. The rightmost sequence traces the Asymptotic Giant Branch (as for field 2). But another sequence, with a constant CO index over 5 magnitude range, is apparent. This traces a different stellar population which corresponds to less evolved, younger stars. The occurrence of relatively recent star formation in the galactic center is already known. However, these observations present an additional diagnostic and allow us to trace with higher accuracy the location and concentration of this star formation event (see Fig. 4), although it is doubtful that this data set can address whether these stars have been formed in situ or fell into the core after their formation.
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Fig. 2 A K‘ image of the central 40” x 40” of the data set. North is up, East to the left. Sgr A* is in the central bright cluster SW of the image. IRS 8 and its bow shock (Rigaut et a1 2003) are north of Sgr A*, West of the high extinction region.
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Fig. 3 CO index (CO - K continuum) vs K continuum for field 2 (right) and field 1 and 2 (left). The two boxes in the left hand panel isolate stellar populations that are only present within N 15" from Sgr A*. Each of the two boxes is associated with a symbol (circle of square) that has been used to mark the objects contained in this box in Fig.4.
Fig. 4 K' image of field I and 2 (20" x 40"). Sgr A* is located by a blue circle. The symbols identify stars of peculiar CO index (see Fig.3).
References Stoke, A., Grebe], E.K., Brandner, W., Figer, D.F. 2003, A&A, in press Figer, D.F. 2002, IAU symposium 212 Yang, Y., Park, H.S., Lee, M.G., Lee, S.-G. 2003, Journal of the Korean Astronomical Society, submitted Rigaut, F., Geballe, T., Roy, J.-R., Draine, B.T. 2003, these proceedings DePoy, D.L., Sellgren, K., Blum, R.D. 2003, BAAS, 198 Figer, D. F. 2003, these proceedings Tanner, A., Ghez, A., Morris, M., Becklin, E. 2003, these proceedings
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Astron. Nachr./AN 324,No. S1,315-319 (2003) / DO1 10.1002/asna.200385073
Stellar Orbits at the Center of the Milky Way N. Mouawad*’, A. Eckart I , S. Pfalzner I , J. Moultaka I , C. Straubmeier I , R. Spurzem *, R. Schodel 3, and T. Ott
’ LPhysikalisches Institut, Universitat zu Koln, Zulpicher Str.77.50937 Koln, Germany Astronomisches Recheninstitut,Monchhofstr. 12-14,69 120 Heidelberg, Germany
’ Max-Planck-Institut fur extraterrestrischePhysik GiessenbachstraRe,85748 Garching, Germany Key words Galactic Center, kinematics and dynamics, inner cusp, central stellar cluster.
Abstract. During the past ten years, measurements of stellar proper motion, and radial velocities (Eckart et al. 2002, Genzel et al. 2000, Ghez et al. 2000) as well as variable X-ray emission (Baganoffet al. 2001) near the center of the Milky Way have convincingly proven the presence of a super-massive3 million solar masses black hole in the center of our Galaxy. We discuss the possible amount of the unresolved mass present at the inner cusp, and its translation into Newtonian periastron-shifts for stellar orbits in the central cluster. For this purpose we use a 4th-order Hermite integrator. Further calculations provide valuable additional information on the three dimensional distribution and dynamics of the He-Stars. We also discuss how future observations with infrared interferometers (LBT, VLTI, Keck) will help to improve our understanding of the dynamics and distribution of the stars in this region,
1
Introduction
As a result of the rapid development of observation techniques, near-infrared (NIR) high resolution imaging and spectroscopic observations with the speckle and adaptive optics (AO) techniques are now able to probe the inner arcsecond of our Galaxy. Observations were made with the MPE speckle camera SHARP at the ESO NTT and the new adaptive optics CONICA/NAOS at the UT4 of the Very Large Telescope (VLT). In the gravitational potential at the center of the Milky Way, the stars show large orbital velocities. In the central arcsecond those can be observed as proper motions via repeated imaging at the highest possible resolution. Ten year of observations have provide sufficient data determine, for the first time, a unique Keplerian orbit for the star S2 (Schodel ct al 2002) and to measure the enclosed dark mass down to a distance of a mere 0.6 mpc from Sgr A*. With this result many alternative scenarios appeared to be unrealistic, leaving a central super-massive black hole as the only plausible explanation. In addition a central stellar cusp was discovered. In the following we discuss how the presence of an extended mass influences the orbits of the inner most stars. The discussion of the central cusp mass is carried out via several steps: First, we calculate by a direct integration the amount of the stellar mass present within the 0.55” radius of Sgr A* due to the cusp (Genzel et al. 2003 submitted). Next, we briefly discuss how the current data can be used to put an upper limit on the cusp mass and the M/L ratio of the cusp. We study the periastron-shift of the S2 orbit, using a 4th order hermit integrator. Finally, we give a first estimate of the line of sight positions of the He I emission line stars, which are enigmatically present in the central cluster at distances not bigger then about 400 mpc from the center. * Corresponding author: e-mail: nelly@phl mi-koeln.de, Phone: +49 221 470 3495. Fax: +49 221 470 5 162
@ 2003 WILEY-VCH Verlag GmbH & Co KGrA. Wcinheim
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2 Mass Estimate of the Inner Cusp 2.1
Stellar Density
With the highest spatial resolution observations presently available in the near-infrared (50 - 60 mas), spatial scales from light hours to a few light years can be probed. The new CONICA/NAOS data were estimated from direct and crowding corrected K, 5 17 and direct ( H S 19) stellar counts in annuli centered on the position of Sgr A*. They clearly confirm the presence of a local stellar concentration, a cusp centered on Sgr A* , indicated earlier by the SHARPNTT and KECK data (Eckart et al. 1995, Alexander 2000). To estimate the cusp mass, we were able to fit the combined SHARP and CONICA stellar count data with a superposition of several Plummer models of the form: p(r)=p(O)[ l+(rR)2]-"/2 (a=S),with different densities p(O), and different core radii. Fig. 1 shows three different fits. The dotted curve gives the fit for the inner cusp with a Plummer model of a core radius R= 0.55" and a spatial density p(r)=4.35x lo7 Ma pcc3. The solid line shows the sum of that initial model with a further inner Plummer model of R= 0.135" and p(r)=6.5x*108 M, p C 3 . The average fit is shown as the large dotted curve. It is similar to the solid curve but with a smaller density of p(r)=3.2Sx 10' Ma pcP3 of the additional R= 0.135" component. Surface density/Distance from SgrA': plummer model fit
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While there is no circular polarization apparent at millimeter wavelengths, we have recently detected linear polarization with the BIMA array at a frequency of 230 GHz (Figure 1). This confirms the JCMT detection at millimeter and submillimeter wavelengths (Aitken et al. 2000). The BIMA observations have an arcsecond beam size, smaller than the JCMT beam by a factor > 100. This allows us to exclude the effects of polarized dust and unpolarized free-free emission with a great degree of certainty. The constant position angle in the upper and lower sideband of the BIMA observations at 230 GHz places a strong upper limit to the rotation measure (RM) of the accretion environment of 2 x lo6 rad m-'. An RM this small excludes a number of models which require large mass accretion rates onto the black Ma y-'. Thus, the low hole. These include ADAF and Bondi-Hoyle models which require luminosity of Sgr A* is due to a low accretion rate rather than to a radiatively inefficient accretion flow.
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Fig. 2 Position angle as a function of frequency. Triangles are the A00 data. Squares are the BIMA data. The solid line is a fit for the RM excluding the A00 230 GHz result. The best fit is -4.3 f 0.1 x lo5 rad m-2 with a zero-wavelength position angle of 181 & 2 degrees.
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We have a tentative result of RM 4 x lo5 rad m-2 based on the BIMA and JCMT measurements. However, these measurements are not fully consistent and re-observation is necessary. An actual measurement of the RM will permit us to exclude even lower accretion rate models. We will also be able to probe changes in the accretion environment as a function of time. 10% at 230 The linear polarization spectrum shows a sharp transition from < 1%at 100 GHz to GHz. Bandwidth depolarization is insufficient to account for this transition while beam depolarization is marginally adequate. Beam depolarization requires a fully turbulent medium with a very small outer scale of turbulence. An alternative scenario involves an unpolarized low frequency component and a highly polarized high frequency component with an inverted spectrum. This is consistent with models that characterize the total intensity spectrum as composed of two separate components. The detection of linear polarization in Sgr A* opens up opportunities to study the immediate environment of a black hole in substantial detail. A number of papers have proposed that general relativistic effects may be detected (Broderick & Blandford, these proceedings; Falcke, Melia & Ago1 2000; Bromley & Melia 2001).
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2 LLAGN The polarization properties of Sgr A* are distinct from those of high powered AGN. In particular, linear polarization dominates circular polarization by typically an order of magnitude in these sources. These powerful AGN are more luminous than Sgr A* by 10 orders of magnitude. To determine whether the
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difference in polarization is an effect of luminosity, we have studied the linear and circular polarization of a sample of nearby LLAGN. In Figure 3 we summarize our observations of 9 LLAGN with the VLA at 8.4 GHz (Bower, Falcke & Mellon 2002). With the exception of M87, all sources showed weak or no linear polarization. Linear
polarization was only detected in two sources at a level of a few tenths of a percent. In NGC 4579 there is a potential detection of RM= 7 x lo4 rad m-’. Circular polarization is only detected convincingly in one source M81*, which we discuss in the next section. The absence of linear polarization for these sources is different from the case of higher luminosity sources and is similar to the case of Sgr A*. The result can be readily explained through Faraday depolarization. An RM- lo5 rad m-’ can produce bandwidth depolarization in these sources at these frequencies. An RM- lo3 rad m-’ can produce beam depolarization. We have demonstrated that both the accretion region and galactic environment of Sgr A* can readily lead to RMs this large. Nevertheless, one cannot exclude the effects of low jet power. If low luminosity jets do not exhibit the same degree of field order as their higher-powered cousins, then they will be weakly polarized. They may lack the powerful shocks that order the magnetic field in high luminosity sources. We note also that there is a clear trend in spectral index and circular polarization strength. Of objects that we have studied, only Sgr A* and M8 1* show circular polarization and only Sgr A* and M8 1 * have inverted spectral indices. This supports the hypothesis that circular polarization is due to an opacity effect in the field. However, only a small number of sources are included in this analysis.
Astron. Nachr./AN 324, No. S1 12003)
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3 MU* We have explored the linear polarization properties of the LLAGN M8 I * in more depth (Brunthaler, Bower & Falcke 2001). These investigations have shown that the polarization continues to exhibit similarity with Sgr A* as we observe at higher frequencies and study the variability properties. The presence of a jet in M81* suggests that the polarization properties of both sources are dominated by a jet. In Figure 4 we show that linear polarization is absent up to a frequency of 22 GHz. This implies lower limits to the RM greater by a factor of N 7 over those for the 8.4 GHz survey. These RM limits are still not surprising given the expected particle densities and magnetic field strengths near the black hole. The circular polarization properties of M81* are also similar to that of Sgr A* (Figure 5). We see that circular polarization is detected at 4.8, 8.4 and 14.9 GHz with a magnitude as high as 1.5%. The degree of variability increases with frequency. The lightcurves suggest episodic activity that is also characteristic in Sgr A*. The highest point in the circular polarization light curve occurs 10 days after a bright fare in the total intensity. As the flare decays, the high frequency circular polarization disappears.
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4
Summary
We have presented here a range of observations of polarization in LLAGN, including Sgr A* and M81*. These sources differ markedly from higher luminosity AGN. These differences are consistent with smallscale, low-power jets which see the high density accretion regions andor galactic HI1 regions. Higher frequency observations may reveal a marked transition in the polarization properties of LLAGN other than
G. C. Bower: Polarization of Sgr A*
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Circular Polarizationin M81 (blue squares) and J1053+704 (black triangles)
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Fig. 5 Circular polarization of M81” and a calibrator as a function of time at three frequencies. The circular polarizaiton properties of M81* are similar to those of Sgr A*. Sgr A* as Faraday effects weaken. These observations are at the edge of capability for current millimeter interferometers. Future arrays such as CARMA and ALMA will be able to systematically survey a broad sample of LLAGN. We will be able to determine the nature of their accretion environments, including the role of advection, convection and outflows. We will b e able to explore the stability of magnetic field structures, the presence of black hole spin and general relativistic effects in the vicinity of the black hole. A similar version of this paper was also submitted to the Proceedings of the Amsterdam Workshop on
Circular Polarization, 2003, J.-I! Macquardt and R. Fendel; eds.
References Aitken, D. K., Greaves, J., Chrysostomou, A., Jenness, T., Holland, W., Hough, J. H., Pierce-Price, D., & Richer, J. 2000, ApJ, 534, L173 Bower, G. c.,Backer, D. C., Zhao, J. H., Goss, M., & Falcke, H, 199ga, A ~ J521, , 582 Baganoff, F., et al., 2003, these proceedings Beckert, T., et al., 2003, these proceedings Bower, G. C., Falcke, H., & Backer, D. C. 1999b. ApJL, 523, L29 Bower, G. C., Falcke, H., Sault, R. J., & Backer, D. C. 2002a, ApJ, 571,843 Bower, G. C., Wright, M. C. H., Backer, D. C., & Falcke, H. 1999c, ApJ, 527,851 Bower, G. C., Wright, M. C. H., Falcke, H., & Backer, D. C. 2001, ApJL, 555, L103 Bower, G. C., Wright, M. C. H., Falcke, H., & Backer, D. C. 2003, ApJ, in press Bromley, B. C., Melia, F., & Liu, S. 2001, ApJL, 555, L83 Brunthaler, A., Bower, G. C., Falcke, H., & Mellon, R. R. 2001, ApJL., 560, L123 Rayner, D.P., Nonis, R.P. & Sault, R.J., 2000, MNRAS, 319,484 Ruszkowski, M. & Begelman, M.C., 2002, ApJ, 573,485
Astron. Nachr./AN 324, No. S1.355-361 (20031 / DO1 10.1002/asna.200385096
Intrinsic Radio Variability of Sgr A* Jun-Hui Zhao*' I Harvard-Smithsonian Center for Astrophysics, 60 Garden St., MS 78, Cambridge, MA 02138 Key words Black hole, radio continuum, accretion disk.
Abstract. We review and summarize the results on the variability of Sgr A* based on the recent monitoring programs with the SMA at 1.3 and 0.87 nun, and the VLA at 2, 1.3 and 0.7 cm. We discuss the flares at 1.3 mm and a cross-correlation of the SMA flux density at 1.3 mm with the VLA data at 1.3 cm. We also present a preliminary result on the double quasi-periodic oscillation (DQPO) in flux density of Sgr A* based on our analysis of radio light curves observed with the VLA at 1.3 cm.
1 Introduction Variations in the radio flux density of Sgr A* have been known for more than two decades (Brown & Lo 1982). The nature of the radio variability in flux density appears to be far more complicated than we had thought. At long wavelengths, the flux density of Sgr A* might be modulated by scintillation due to the turbulence in the ISM (Zhao et al. 1989). The early VLA monitoring campaign at 20, 6, 3.6, 2 and 1.3 cm during the period 1990-1993 shows that the fractional amplitude variations increased towards short wavelengths and that the rate of radio flares appeared to he about three per year, suggesting that the intrinsic radio outburst occurs in Sgr A* (Zhao et al., 1992; and Zhao & Goss 1993). The typical time scale of these radio flares is about a month. The observed large amplitude variations in flux densities at 3 mm (Wright & Backer 1993; Tsuboi, Miyazaki & Tsutsumi 1999) are consistent with the wavelength-dependence of the variability as observed at centimeter wavelengths. The monitoring observations at 1.3 mm with the partially completed SMA (Moran 1998) also suggest large amplitude flares. The presence of a 106f10day cycle in the radio variability of Sgr A* was suggested from an analysis of data observed with the VLA in the period of 1977-1999 (Zhao, Bower and Goss 2001). Similar Huctuation in flux density was also observed in the Green Bank Interferometer monitoring data (Falcke, 1999). In this paper, we summarize the result from recent 25 epochs of observations with the partially completed SMA and discuss the preliminary results from a power spectrum density (PSD) analysis of the radio light curve observed with the VLA at 1.3 cm.
2 Flares at 1 mm Observarions of Sgr A* at I .3 mm and 0.87 mm were made using the partially completed SMA with three or four antennas and baselines ranging from 7 to 55 kilo wavelengths at 1.3 mm. A complex, large scale structure in the central region is automatically filtered away. For baselines 20 kX and longer, Sgr A* appears (0 he the dominant source and the confusing flux density from the surrounding medium appears to be less significant at these two short wavelengths. Fig. 1 shows the SMA light curve at 1.3 mm suggesting that Sgi- .A':: varies significantly. Three possible Hares were observed during the monitoring period of 15 I I N - ~ I ~{' ~. ~C. CFig. 1 and Zhao et a1 2003). The time scale of the variation of a month is consistent with
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* Corresponding author: e-mail: jzhaoQcfa.harvard.edu,Phone: tl617 496 7895, Fax: +1617 496 7554
@ 2003 WILEY-VCH Verlag GmbH & C o KGaA, Weinhem
Jun-Hui Zhao: Radio Variability of Sgr A*
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Fig. 1 The SMA light curve of Sgr A* observed at 1.3 mm in the period between March 2001 and May 2002
2.1 Spectra and Sub-millimeter Component Spectra during a flare and a minimum are plotted (see Fig. 2). The spectral index a ( S , 0: va) appears to be 0.1+0.1 at 100 GHz and below, and 1.5?::: between 232 and 345 GHz, suggesting a break frequency in spectral index at 100 GHz or higher. A flux density excess towards sub-millimeter wavelengths has been observed (Zylka, Mezger, & Lesch 1992; Serabyn et al. 1997; Falcke et al. 1998). The overall spectrum 1(v/v~)~'+S ~ ( Y / V ~Three ) ~ ' .sets of combination of crl can fit two power-law components, i.e., S, = S and a 2 are used in the fitting. First, for c u l = 0 and a2 = 2 (dashed lines in Fig. 2), the sub-millimeter component corresponds to the thermal synchrotron emission either arising from the inner region of the accretion disk (Liu & Melia 2002) or produced from a jet-nozzle (Falcke & Markoff 2000). Secondly, a1 = 0 and a2 = 2.5 (solid lines in Fig. 2), the spectral index of 2.5 suggests that a homogeneous opaque, nonthermal synchrotron source might be present in the inner region of the accretion flow. Such a model appears to be plausible if one considers the non-thermal synchrotron particles to be accelerated inside the compact source, perhaps within a jet nozzle as has been proposed for the case of NGC 4258 (Yuan et al. 2002). Sl(v/v1)O 25ezp(-v/vo), Finally, if the low frequency component has an exponential cut-off, i.e. at vo 7 5 GHz, a smaller value of a 2 1.5 (dashed-dotted lines in Fig. 2 ) for the sub-millimeter component is also consistent with a spectrum produced from the ADAF model. A gradient of T, in the ADAF depresses the rising part of the spectrum (Narayan et al., 1998). The observed spectrum suggests an opaque nature of the sub-millimeter component at 1.3 mm and perhaps also 0.87 mm. Observations at shorter sub-millimeter wavelengths appear to be critical to differentiate between the models. The spectrum in a minimum state is also shown here. The excess at 1.3 mm appears to be less significant. N
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2.2 Day-to-Day Variability and Intra-day Variability Based on the sparse data, marginal day-to-day variations at a level 2-3u (or 20-30%) were observed during Flare 1 and Flare 2 as well as May 2002. Intra-day variations on short time scales were searched based on the 24 epochs of observations at 1.3 mm (total time of 100 hr). No evidence for significant variations on a time scale of 1 hr has been found, N
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Correlation with the VLA data
Fig. 3 shows the SMA light curve at 1.3 mm overlaid on the VLA light curve at 1.3 cm. The SMA data appears to show a correlation with the light curves observed with the VLA. A quantitative analysis of cross-correlation between the light curves at 1.3 mm and 1.3 cm suggests a global delay of &lay > 3d (Zhao et a]. 2003). The global delay is consistent with a model consisting of a flare that occurs from inside out starting from short wavelengths and then continuing to longer wavelengths. Assuming a global delay of &lay > 3d and a source size of 40 R,3,, an expansion velocity, vezp -1200 km spl or < 0.004 c, is inferred, which is far below the escape velocity of 0.1 c at r ,-..40 Rsr. The bulk kinetic energy associated with the flares appears to be too small to power a noticeable collimated jet in Sgr A* during the SMA monitoring period. The inferred small expansion velocity may imply that other processes contribute to the transport of high-energy particles; e.g. diffusion and convection may also play a role in powering Sgr A” at lower radio frequencies.
Jun-Hui Zhao: Radio Variabilitv of Spr A*
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In Fig. 3, we also mark the X-ray flares observed in Oct 2000 (Baganoff et al. 2001) and May 2002 during the multi-wavelength campaign (Baganoff et al. 2003). The first observed X-ray flare appcars to be 10 d prior to the peak of the quasi-periodic fluctuation observed at I .3 cm. Corresponding to the multiple X-ray flares observed during the last week of May 2002, the flux density measured with SMA at 1.3 mm shows a 20 increase. Unfortunately, no further SMA observations were made following the multiwavelength campaign of May 2002. However, a 2 0 peak was observed with the VLA at 1.3 cm about 4 weeks after the multiple X-ray flares. The observed properties for the flares at 1.3 mm (time scale of a month and amplitude fluctuation of a factor of a few) are different from those of the X-ray flares (time scale of 1 hr and amplitude fluctuation of a factor of 10 or larger). The lack of strong flares on short time scale at 1 mm places a critical constraint on the models of the inverse Compton scattering as has been proposed for the short duration X-ray flares. Considering the opaque nature of the sub-millimeter component at 1.3 mm, the X-ray flares could remain hidden at 1.3 mm due to self-absorption. ",
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Fig. 3 The SMA light curve at 1.3 mm (squares) is overkdid with the VLA light curve at 1.3 cm (dots). The solid curve is a smoothed 1.3 cm light curve. The longer period component (P2 333 d) of the DQPOs is noticeable. The X-ray flares observed with Chandra in Oct. 2000 and May 2002 are marked.
3 Variability at 1 Centimeter 3.1
New VLA Light Curves
Fig. 4 shows the new weekly sampled radio light curves observed with the VLA at 2, 1.3 and 0.7 cm during the period between June 2000 to the end of 2002 (Herrnstein et al. 2003 in preparation). The spectral indices of a1 3cm/2cm and a0 7cm/l :irm are also calculated. The mean values are a1 .3cm/2cm = 0.3 f0.2 and ~ ~ 0 , 7 3c7rL ~ ~= / 0.12 1 0.19.
*
3.2
Double Quasi-Periodic Oscillation
Fig. 5 shows the light curve at 1.3 cm obtained from the VLA archival data over the period 1990-1999 (upper panel; Zhao, Bower & Goss 2001) with a total number of 92 data points and a mean sampling interval At=39d. The PSD (middle panel) is derived with the Lomb Algorithm from a portion of the light curve observed at 1.3 cm during 1990 March to 1993 October over a 1400-d period with At=22d. We note that for all the PSD plots, no baseline subtraction was applied to the data prior to the application of the Lomb algorithm. In addition to the strong feature at near O.O1d-l (NlOOd),there is a weak periodic signal at 0.004d-' (-240d). The PSD (bottom panel) confirmed the presence of the double oscillation signals at 0.01d-' and 0.004dC1 from all the data observed from 1990 to 1999 at 1.3cm. The double oscillation signals appear to be also present in the densely sampled light curves observed over an interval from June 2000 to December 2002 at 1.3 cm (see Fig. 3) with a mean sampling interval
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of 8d. All three PSDs derived from the different observing periods during the course of the monitoring program are dominated by two features at ,-- 0.0075 d-' (133 d) and ,-- 0.003dP1 (333 d), while the power of the shorter period is decreasing in magnitude (see Fig. 6). The high power oscillation feature of P2333 d is also shown as a sinusoidal oscillation in a smoothed light curve ( nearly three completed cycles over 9 14 d, see Fig. 3). The weaker, minor signals (- 2 0 peaks) appear to correspond to P1 133 d mode, which is superimposed on the P2 oscillation mode. The SMA flares at 1.3 mm are likely associated with the PI oscillation mode. N
4
Summary
With nearly 200 epochs over the past 13 years, the variability of Sgr A* at I .3 cm is dominated by two oscillation modes with periods of P1 = l0Od - 133d and P 2 = 250d - 333d based on both white noise and l/f noise analysis. Both the intensity and the period of the oscillation drift in time, but the ratio of the periods maintains a nearly constant ratio of 2.4 -2.5. The DQPO (double quasi-periodic oscillation)
Jun-Hui Zhao: Radio Variability of Sgr A*
360
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may well provide crucial information about the activity near the supermassive black hole, which may be triggered by an orbital resonance occurring in the central few thousand Schwartzschild radius (Zhao et al. 2003, in preparation). The SMA flares at 1.3 mm appear to correspond to the P1 oscillation mode revealed from the VLA light curve observed at 1.3 cm, suggesting that Sgr A* may indeed be regularly powered by the activity near the supermassive black hole. The flares at sub-millimeter wavelengths might be a result of collective mass ejections associated with the X-ray flares that originate from the region near the event horizon where the accretion flow is falling into the black hole. The emitting electrons at sub-millimeter and longer radio wavelengths are likely re-processed from the X-flaring plasma through particle acceleration processes, such as shock and magnetic field reconnection. Sgr A* appears to be opaque at 1.3 mm and perhaps at 0.87 mm. Higher frequency observations of Sgr A* at sub-millimeter wavelengths, where the sub-millimeter component may become partially optically thin, are crucial to our understanding of the nature of the flares observed at centimeter, millimeter/sub-millimeter and X-ray wavelengths.
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Fig. 6 The PSDs derived from the new VLA observations at 1.3 cm during 2000-6-21 to 2002-9-10 (see Fig. 1 for the light curve). We break the light curve into three portions with different number of measurements covering different observing periods in order to show the double oscillation features and their drifting behavior in both intensity and frequency as a function of time. Fig. 3 shows the PSDs derived from three different time periods: 1 ) 2000-6-21 to 2001-7-9 ( S O data points), 2) 2000-6-21 to 2002-2-18 (75 data points), 3) 2000-6-21 to 2002-9-10 (100 data points). Acknowledgements The author especially thanks Miller Goss for his encouragement, interest, and many stimulating discussions during the course of this long term collaboration project. This review is based on collaborations with Goss, Bower, Herrnstein, Ho, Lo, Pegg, Tsutsumi and Young. The author is grateful to Jim Moran, Director of the SMA, for encouragement, support and helpful discussion. The author also thanks Barry Clark for scheduling the VLA monitoring program. The Very Large Array (VLA) is operated by the National Radio Astronomy Observatory (NRAO). The NRAO is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.
References Baganoff, F. K., Bautz, M. W., Brandt, W. N. et al. 2001, Nature, 413,45 Baganoff, F. K. et al. 2003, these proceedings Brown, R. L. and Lo, K. Y. 1982, ApJ, 253, 108 Matsuo, H., Teuben, P., Zhao, J.-H., & Zylka, R. 1998, ApJ, 499, 73 1 Falcke, H., & Markoff, S. 2000, AA, 362, 113 Liu, S. and Melia, F. 2002, ApJL, 566, L77 Moran, J. M. 1998, in: Proc. SPIE, Vol. 3357: Advanced Technology MMW, Radio & Terahertz Telescopes, ed. Thomas, G. Phillips, p. 208 Narayan, R., Mahadevan, R., Grindly, J., Popham, R., & Cammie, C. 1998, ApJ, 492,554 Press, W. H., et al. 1992, Numerical Recipes in C, The Art of Scientific Computing, Second Edition, Cambridge University Press Serabyn, E., Carlstrom, J., Lay. O., Lis, D., Hunter, T., & Lacy, J. 1997, ApJL, 490, L77 Tsuboi, M., Miyazaki, A. & Tsutsumi, T. 1999, ASP Conf. Series 186, p105 & Tsuboi, M. 2002, BAAS, #44.09 Wright, M. and Backer, D. C. 1993, ApJ, 417,560 Yuan, F., Markoff, S., & Falcke, H., Biermann, P. L. 2002, AA, 391, 139 Zhao, J.-H., Ekers, R.D., Goss, W.M., Lo, K.Y. & Narayan R. 1989, IAU Symp. 136,535 Zhao, J.-H., Goss, W. M., Lo, K. Y. and Ekers, R. D. 1992, ASP Conf. Series 31,295 Zhao, J:H. and Goss, W. M. 1993, in: Sub-arcsecond Radio Astronomy, R.J. Davis and R. S. Booth, Cambridge University Press, 38 Zhao, J.-H., Bower, G. C., Goss, W. M. 2001, ApJ, 547, L29 Zhao, J.-H., Young, K. H., Herrnstein, R. M., Ho, P.T.P., Tsutsumi, T., Lo, K. Y., Goss, W. M., & Bower, G. C. 2003, ApJL, 586, March 20 issue, in press Zylka, R., Mezger, P. G., & Lesch, H. 1992, AA, 261, 119
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Astron. Nachr./AN 324, No. S1,363-369 (2003) / DO1 10.1002/asna.200385071
Flares of Sagittarius A* at Short Millimeter Wavelengths Atsushi Miyazaki*',Takahiro Tsutsumi **
2,
and Masato Tsuboi ***
' Nobeyama Radio Observatoryt, National Astronomical Observatory of Japan, Minamimaki, Minamisaku, Nagano 384-1305, Japan National Astronomical Observatory of Japan, 2-21-1 Osawa, Mitaka, Tokyo 181-8588,Japan Institute of Astrophysics and Planetary Science, lbaraki University, 2-1-1 Bunkyou, Mito, Ibaraki 3108512, Japan
Key words galaxies: nuclei, Galaxy: center, radio continuum Abstract. We have performed intensity monitoring observations toward the Galactic center compact nonthermal radio source, Sagittarius A*, at A= 3 and 2 mm (100 and 140 GHz) from 1996 to 2002 using the Nobeyama Millimeter Array (NMA). In 1996, 1997, 1998, and 2000, the observations were performed during a period of one to two months for each year in a single (inteiinediate resolution) array configuration in order to exclude a systematic error caused by sampling different source structures with the different beam sizes. From November 2000 to May 2001, and from November 2001 to May 2002, the observations were carried out over several months in order to investigate the longer term variability. We have detected two flares in March 2000 and April 2002 as well as a flare observed in March 1998 (Tsuboi et al. 1999). The peak flux densities of the flares were 2-3 Jy at 3 mm, while the mean quiescent flux densities was -1 Jy. In particular for the March 2000 flare, the flux densities of Sgr A'at 2 mm had also reached a peak, -4 Jy, on 8 March and increased AS/S- 300%. The flux density decreased by half in a day. We have detected the intra-day variability of Sgr A*. The upper limit for size of the variable component estimated from timescale of the flare is a few tens of AU. For spectra made from our Observations at mm-wavelengths, the variability of flux density increases with frequency. It appears that the variability in the flare propagates from higher to lower frequency. The folded lightcurve with 106 days cycle, which determined from the analysis of the VLA cm-wavelength data, shows distinct high and low activity states. These results provide evidence for quasiperiodic variability of the flux density of Sgr A*at mm-wavelengths.
1 Introduction Sagittarius A* (Sgr A*) is a unique compact nonthermal radio source located at the dynamical center of the Galaxy (e.g. Eckart & Genzel 1996). Sgr A*is suggested to be associated with a supermassive black hole, with about 2.6 x lo6 Ma, at the Galactic center (e.g. Eckart & Genzel 1996; Ghez et al. 1998). The radio emission of Sgr A*is thought to be powered by the gravitational potential energy released by matter as it accretes onto a supermassive black hole. However, the true radiative nature is still not well understood. Because this source is embedded in thick thermal material, it is practically difficult to observe its fine structures with the present VLBI (very long baseline interferometry; Doeleman et al. 2001). Time variability observation is a powerful and alternate tool to reveal the emission mechanism and the structure of Sgr A*. If the time variability is intrinsic to the source, it should be tightly related to the emission * Corresponding author: e-mail:
[email protected],Phone: +81267 984381, Fax: +81267 982923 ** Corresponding author: e-mail:
[email protected] *** Corresponding author: e-mail: tsuboi @mx.ibaraki.ac.jp Noheyama Radio Observatory (NRO) is a branch of the National Astronomical Observatory, an inter-university research institute operated by the Ministry of Education, Culture, Sports, Science and Technology.
@ 2003 WILEY-VCH Verlag GmbH & Co. KGaA. Weinheirn
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mechanism and the structure of the emitting region. If the origin is not intrinsic, the observation should at least provide the information about thermal materials around Sgr A*. Time Variability in the flux density of Sgr A*at centimeter wavelengths has been studied in the last decade (e.g. Zhao et al. 1992). The variability below 10 GHz is dominated by refractive interstellar scintillation. However, Zhao et al. (2001) reported the presence of a 106 day cycle of variability at centimeter wavelengths based on an analysis of data observed with the Very Large Array over the past 20 years. Wright & Backer (1993) reported that significant flux variations at X=3 mm occurred in a month during the decay of a flare at short centimeter wavelengths in 1990. The millimeter variability is believed to reflect the intrinsic activity of Sgr A*. However, the millimeter variability is not well established because of the lack of systematic monitoring observations. Thus it is important to carry out systematic and multi-epoch monitoring observations at millimeter wavelengths in order to probe the intrinsic variability of Sgr A*. Tsuboi, Miyazaki, & Tsutsumi (1999) had reported the flare of Sgr A’at 3 mm in March of 1998 (Tsuboi et al. 1999; Miyazaki et al. 1999). The flux density of Sgr A*at 100 GHz was flared over AS/S=lOO% in a week and decreased to the mean flux density within two weeks. The flux density at 140 GHz during the flare also increased by more than 100% (AS/S). On the other hand, Zhao et al. (2003) have also detected flares of Sgr A*at 1.3 mm wavelength using the Suhmillimeter Array (SMA). Recently, Baganoff et al. (2001) detected an X-ray flare lasting about 10 ks and with a peak luminosity -50 times higher than the quiescent state by Chandra observations. The relation between X-ray and radio flares is interesting. Sgr A*is a relatively weak (- 1 Jy) compact component embedded in the extended and strong HI1 region of Sgr A West. The contribution of the free-free emission from the extended structure is significant even at millimeter wavelengths. It is necessary to observe with higher angular resolution ( 5 a few arcsec.) to discriminate the compact component from the extended components. We have conducted intensity monitoring experiments toward Sgr A*at millimeter wavelengths using the Nobeyama Millimeter Array (NMA) at the Nobeyama Radio Observatory (NRO).
2 Observations & Calibrations We have performed intensity monitoring observations toward Sgr A*at 100 and 140 GHz band (X=3 and 2 mm) using the NMA, a six-element interferometer at the NRO from 1996 to 2002. In 1996, 1997, 1998, and 2000 the observations were performed during a period of one to two months for each year. Different configurations of the antennas could introduce a bias in the result caused by sampling different source structures with the different beam sizes. While the highest angular resolution can be achieved by the array configuration with the longest baseline, it always requires the best weather conditions and it is not suitable for intensity monitoring. Therefore, we used a single (intermediate resolution) array configuration of the NMA, “C-configuration”, in these period in order to exclude a systematic error as above instance. Each epoch consists of a set of sequent observations of 2-3 days. The epochs of the observations were separated by about 5 days and carried on during a period of one to two months. From November 2000 to May 2001, and from November 2001 to May 2002, the observations were carried out over several months using the various array configurations, which included “AB-configuration” (long baseline), and “D-configuration” (short baseline), in order to investigate the longer term variability as a known variability at centimeter wavelengths. Each epoch consists of a set of sequent observations of 2 days. The epochs were separated by 10 days to 2 weeks. We used double side band (DSB) SIS receivers at the 100 and 140 GHz bands as the front-ends. The data in 1996 were acquired at observing frequencies of 102 and 146 GHz using the FX correlator with 320 MHz bandwidth. The data from 1997 to 2002 were acquired using the Ultra Wide Band Correlator (UWBC; Okumura et al. 2000) with 1 GHz bandwidth. The UWBC is able to observe in the lower and upper side hands separated by 12 GHz simultaneously. The observed frequencies form 1997 to 2002 were 90 and 102 GHz for the 3 mm band and 134 and 146 GHz for the 2 mm band. Because the observations for
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Sgr A’at 2 mm band requires the best weather conditions and the phase stability, these observations were made less frequently. The instrumental gain and phase were calibrated by alternating observations of Sgr A*and NRAO 530 at about 20 min. intervals. We also observed an additional phase calibrator, 1830-210 of the known QSO, from 2000 to 2002. The flux density scale was established from the observations of Uranus or Neptune. The absolute flux density scales are typical accurate to about 15% for the 100 GHz band and about 20% for the 140 GHz band. To check the validity of flux calibration, Sgr B2(M) is also observed in synthesis mode. We assumed the flux density of Sgr B2(M), which is HIT region, to be constant. Sgr A*
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C-configuration of NMA contains U-V distances at 100 and 140 GHz bands of -7-55 kX and -1077 kX, respectively. AB- and D-configuration contains U-V distances at 100 GHz band of -20-100 kX and -4-25 kX, respectively. We made spatially filtered maps from the U-V data by restricting to spatial 2 2 5 k X ) in order to suppress the contamination from the frequencies larger than 25 kX ( ( U z V2)’/’ surrounding extended components around Sgr A*(for the data taken with the compact array configuration, D-configuration, 17kX). We use the peak flux density on the maps to look for variability. Figure 1 shows the spatially filtered maps at 102 GHz in 7 March 2000, and at 146 GHz in 8 March 2000 which are made from the (I-V data by restricting to spatial frequencies larger than 25 kX. The typical size of the synthesized beam with uniform weighting for C-configuration is 3” x 6” at 100 GHz and 2” x 4” at 140 GHz. The resultant peak flux densities are corrected for the effect that the flux density is reduced due to phase errors by atmospheric fluctuations (The fractions of the reduction are 10-20% for 100 GHz and 20-40% for 140 GHz).
+
3 Results & Discussion Figure 2a shows the light curve of Sgr A*at 100 GHz band constructed from all monitor data from 1996 to 2002. Total number of observations is about 60 days. The light curve shows that Sgr A*has distinct active and quiescent phases. There are three active phases in March 1998, March 2000, and April 2002 (Flare I, Flare 11, and Flare 111, respectively, indicated by the arrows in Figure 2a) with duration of, roughly, one month. Mean flux densities of Sgr A*in quiescent phase except for the flare phase are 1.1=k 0.2 Jy and 1.2 k 0.2 Jy at 90 and 102 GHz, respectively. These flare phases are account for 20 % of all data. The increase (AS/S) of the flux density during the flare reached to 200%. The flare in March 1998 (Flare I) has been reported in the previous papers (Tsuboi et al. 1999; Miyazaki et al. 1999). Baganoff et al. (2001) N
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Fig. 2 The light curve of Sgr A'at 100 GHz band from 1996 to 2002. The open and filled circles show flux densities of Sgr A'at 90 and 102 GHz, respectively. (a) The light curve of whole flux density data from 1996 to 2002. There are three active phases in March 1998, March 2000, and April 2002 (Flare I, Flare 11, and Flare EI,respectively, indicated by the arrows). (b) The light curve in 1996, (c) in 1998, (d) in 2000, and (e) from November 2000 to May 2001. The light curve is separated into the active (in 1998, 2000) and the non-active (in 1996,2000-01) phases. The dash line indicates the mean flux density at 100 GHz band.
had detected the X-ray flare of Sgr A*with Chanrlra X-ray Observatory at 26-27 October 2000. However, there are no observations by us around the end of October 2000 (Figure 2e) to make any comparison. In 22 November 2000, our observation detected a flux density of 1.0 Jy at 100 GHz band. N
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3.1 Variability Feature in the Flares In the March 2000 Hare (Flare II), the Hux densities of Sgr A*at 3 mm had reached 2 peaks (-3 Jy) on 7 and 21 March and increased AS/Sw 200% for both. The Flare I1 was also observed at 140 GHz band. Figure 3a shows the light curves of Sgr A*at 100 and 140 GHz bands in March 2000. There were several flares in the active phase, most notably a steep peak at 140 GHz band observed on 8 March 2000. The flux densities of Sgr A*at the peak were 3.5 f 0.7 Jy and 3.9 f0.8 Jy, at 134 and 146 GHz, respectively, while the averaged quiescent Hux was about 1 Jy at 140 GHz band. The Hux density then decreased to 2.2 3~ 0.4 Jy at 146 GHz on the subsequent day, 9 March 2000. The half decay timescale of the flare at 146 GHz, was at most 24 hours. The flare amplitude from the mean flux density level (AS/S) reached to about 300% at 146 GHz, which is larger than that at 100 GHz band (- 200%), indicating the variability increases with frequency. We measured the decay timescale of the Hare to be one day (8-9 March 2000). The upper limit of the timescale estimated from time to increase by 100% is about 12 hours. Figure 3b shows the light curves of Hare of Sgr A*at 146 GHz in 8 March 2000. We divided the data set observed in several parts and measured the flux density of Sgr A'from the divided data. The flux density of Sgr A*averaged over 5 to 15 minutes bin at 146 GHz increased from 3.5 to 4.7 Jy between 15h45m UT to 16h 15m UT in 8 March 2000 (Miyazaki et al. 2003). The relative error within one observation session was estimated to be about 5% at 140 GHz band. Thus the 30% Hux increase in 30 minutes is probably real. The time scale that the Hux density increased by 100% is estimated to be about 1.5 hours assuming that the increase has a constant gradient. The increasing time scale, 1.5 hours, provides the physical size (light crossing size) of the Hare region in accretion disk is compact at or below 10 AU (= 200 R,; Schwarzschild radius, R, = 2GM/c2). (Miyazaki et al. 2003). Moreover, there appear also a rapid increase on Flare I11 in April 2002. The Hux density of 1.4 0.2 Jy in 5 April 2002 at 102 GHz had been doubled in only one day and the Hux density was 2.2 0.3 Jy in 6 April. Similar short-term variation at mm-wavelengths has been reported for the galactic nucleus ofM81 (Sakamoto et al. 2001). N
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Date Fig. 3 (a) The light curve of Sgr A*at I 0 0 and 140 GHz bands in the observation season of 2000. The flux density at 100 GHz was violently changing (circles).There is a steep peak at 140 GHz band at 8 March 2000 (squares). (b) The light curve of Sgr A*at 146 GHz in the observation on 8 March 2000. This figure shows an intra-day variability of the flux density of Sgr A*.
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3.2 Spectrum of Sgr A*and Long Term Variability Figure 4 shows the spectra of Sgr A'at mm-wavelengths (90, 102, 134, and 146 GHz) in several epochs. The observations of each band (100 and 140 GHz band) have been performed at two consecutive days, respectively. The spectral index of Sgr A*at cm-mm wavelength in the quiescent phase are consistent with 0.3, from the extrapolation of cm-wavelengths ( e g Falcke et al. 1998; the expected spectral index, Krichbaum et al. 1997). We did not find the short-mdsubmm excess in the spectrum of Sgr A*over the expected flux determined from the extrapolation of lower frequency data in the quiescent phase. The excess is presumably caused by the variable component in mm-wavelengths (Tsuboi et al. 1999). Moreover, as a conjecture from the light curve of the Flare I1 (figure 3), the variability of flux density increases with frequency. It appears that the (time sequence of) flares propagate from higher to lower frequency. N
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Frequency (GHz) Fig. 4 Spectra of Sgr A*at mm-wavelengths (90, 102, 134, and 146 GHz). The straight lines indicate the best-fitting power law to the flare and quiescent phases in 1998 compiled the data at 2.25 GHz and 8.3 GHz are from the archival data of the Green Bank Interferometer (GBI) of NRAO. The cm-mm spectra were explained by the power law with the spectral index of 0.55 at time of the flare and 0.34 during the quiescent phase (Tsuboi et al. 1999).
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Date Fig. 5 The folded lightcurve of the flux density of Sgr A'at 90 and 102 GHz with a 106 days cycle. The open and filled circles show flux densities of Sgr A'at 90 and 102 GHz, respectively. The dash line indicates the mean flux density at 100 GHz band.
Zhao, Bower and Goss (2001) had reported the presence of a quasi-period of 106 days cycle in the variability at centimeter wavelength of Sgr A*based on an analysis of data observed with the VLA over the 20 years. We folded the our NMA light curve with the 106 days determined from the analysis of the VLA centimeter wavelength data. Figure 5 shows the folded lightcurve of the flux density of Sgr A*at 90 and 102 GHz. In Figure 2, the light curve can be separated into the active phase (flares in 1998 (c), 2000 (d)) and the non-active phase (quiescent in 1996 (b), 2000-01 (e)). The folded lightcurve in Figure 5 shows distinct high and low activity states. These results provide evidence for a quasiperiodic variability of the flux density of Sgr A*at mm-wavelengths.
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We have performed monitoring observations of the flux density of Sgr A*at 3 mm (100 GHz) and 2 mm (140 GHz) bands using the Nobeyama Millimeter Array (NMA) in 1996, 1997,2000,2000-01, and 200102. 1) We detected three flares from Sgr A*in March 1998, March 2000, and April 2002, at 100 and 140 GHz bands. The peak flux densities of the flares were 2-3 Jy at 100 GHz, while the mean quiescent flux densities were -1 Jy. At 140GHz, the flux density increases up to 200% (=AS/S). 2 ) For the March 2000 flare, the flux density of Sgr A*at 2 mm also reached a peak -4 Jy on 8 March and increased AS/S- 300%. The flux was decreased by half in a day. Moreover, we detected the 30% flux increase in 30 minutes on 8 March 2000. This increase demonstrates the intra-day variability of Sgr A*. The upper limit for the size of the variable coniponent estimated from the timescale of the flare is a few tens of AU. 3) For spectra made from our observations at mm-wavelengths, the variability of flux density increases with frequency. It appears that the variability in the flare propagates from higher to lower frequency. 4) The folded lightcurve with the 106 day cycle, determined from the analysis of the VLA cm-wavelength data, shows distinct high and low activity states. These results provide evidence for quasiperiodic variability in the flux density of Sgr A*at mm-wavelengths. Acknowledgements We thdnk the staff of Nh4A group of the Nobeyama Radio Observatory for support in the observation. They also thank T. Kawabata at Bisei Observatory for useful discussions.
References Baganoff, F.K., Bautz, M.W., Brandt, W.N., et al. 2001, Nature, 413, 4.5 Doeleman, S.S., et al. 2001, AJ, 121, 2610 Eckart, A. & Genzel, R. 1996, Nature, 383,415 Falcke, H., et al. 1998, ApJ, 499, 731 Ghez, A.M., Klein, B.L., Moms, M., & Becklin, E.E. 1998, ApJ, 509, 678 Krichbaum, T.P., et al. 1998, A&A, 335, L106 Miyazaki, A.,Tsutsumi, T., & Tsuboi, M., 1999, Advances in Space Research, 23, 977 Miyazaki, A.,Tsutsumi, T., & Tsuboi, M., 2003, in preparation OkUmurd, S.K., et al. 2000, PASJ, 52, 393 Sakamoto, K.,Fukuda, H., Wada, K.,Habe, A. 2001, AJ, 122, 1319 Tsuboi, M., Miyazaki, A., & Tsutsumi, T. 1999, in ASP Conf. Ser. 186, The Central Parsecs of the Galaxy, eds. H. Falcke, A. Cotera, W.J. Duschl, F. Melia, & M.J. Rieke (San Francisco: ASP), p. 105 Wright, M.C.H., & Backer, D.C. 1993, ApJ, 417,560 Zhdo, J.-H., Goss, W.M., Lo, K.Y., & Ekers, R.D. 1992, in Relationships between Active Galactic Nuclei and Starburst Galaxies, ed. A.V. Flilpenko (San Francisco: ASP), p. 295 Zhao, J.-H., Bower, G.C., & Goss, W.M. 2001, ApJL, 547, L29 Zhao, J.-H., et al. 2003, ApJ Letters, in press
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Astron. Nachr./AN 324,No. S1, 371-376 (2003) / DO1 10.1002/asna.200385062
Limits on the Short Term Variability of Sagittarius A* in the Near-Infrared S. D. Hornstein*',A. M. Ghez's2, A. Tanner', M. Morris', and E. E. Becklin' I
Department of Physics and Astronomy, University of California at Los Angeles, Los Angeles, CA 900951562 Institute for Geophysics and Planetary Physics, University of California, Los Angeles, CA 90095-1567
Key words accretion, accretion disks - black hole physics - Galaxy: center - infrared: galaxies X-rays: galaxies - galaxies: jets
Abstract. The detection of X-ray flares by the Chandra X-ray Observatory and XMM-Newtonhas raised the possibility of enhanced emission over a broad range of wavelengths from Sagittarius A*, the suspected 2.6 x 10' M a black hole at the Galactic center, during a flaring event. We have, therefore, reconstructed 3-4 hr data sets from 2 pm speckle and adaptive optics images (Q,,,, = 50 - 100 mas) obtained with the W. M. Keck 10 m telescopes between 1995 and 2002. The results for 25 of these observations were reported by Hornstein et al. (2002) and an additional 11 observations are presented here. In the 36 separate observations, no evidence of any significant excess emission associated with Sgr A* was detected. The lowest of our detection limits gives an observed limit for the quiescent state of Sgr A* of 0.09*0.005 mJy, or, equivalently, a dereddened value of 2.0f0. I rnJy. Under the assumption that there are random 3 hr flares producing both enhanced X-ray and near-infrared emission, our highest limit constrains the variable state of Sgr A* to 20.8 mJy (observed) or 19 mJy (dereddened). These results suggest that the early model favored by Markoff et al. (2002), in which the flare is produced through local heating of relativistic particles surrounding Sgr A* (e.g., a sudden magnetic reconnection event), is unlikely because it predicts peak 2 prn emission of -300 mJy, well above our detection limit.
1 Introduction The variability of Sagittarius A* at X-ray wavelengths (Baganoff et al., 2001, 2003; Goldwurm et a]., 2003) has bolstered the case for associating this source with the suspected 2.6 x lo6 M o black hole at the center of our Galaxy (Eckart & Genzel, 1997; Genzel et al., 1997,2000; Ghez et al., 1998,2000). The first evidence of X-ray variability was detected by Chundru, during which Sgr A* was seen to flare in intensity over a time scale of -3 hr (Baganoff et al., 2001). While the flare's short duration implied a small region of origin, 5 4 0 0 R, (where R, is the Schwarzschildradius = 2GM./c2), its large amplitude, a factor of 50, has raised the possibility of detecting corresponding intensity enhancements at wavelengths outside the X-ray regime. Later observations with XMM-Newton (Goldwurm et al., 2003) as well as follow-up observations by Chundru (Baganoff et al., 2003) both show similar flaring activity. Existing models for Sgr A*'s flared state make very disparate predictions for the emission at wavelengths between the X-ray and radio regimes (Markoff et al., 2001; Liu & Melia, 2002; Narayan, 2002). The wide differences between these models are a result of assuming different geometries (disk vs. jet) and emission mechanisms for the flaring process (e.g., enhanced accretion rates vs. magnetic reconnection). In some models, the predicted emission in the flared state, at infrared wavelengths, dramatically exceeds that of existing detection limits (Genzel & Eckart, 1999; Stolovy et al., 1999; Morris et al., 2001). For example, * Corresponding author: e-mail:
[email protected],Phone: 310-825-4434,Fax: 310-206-2096
@ 2003 WlLEY-VCH Verlag OmbH & Co. KGaA, Weinhem
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the preferred model of Markoff et al. (2001) predicts an observed 2 p m flux density of 13 mJy, or a dereddened flux density of -300 mJy, during the flared state. Unlike the situation at radio wavelengths, where long-term monitoring campaigns have been used to constrain the flared state of Sgr A* (Zhao, Bower, & Goss, 2001), the limited time coverage and spatial resolution of published IR experiments prevent meaningful constraints on the flared state’s IR emission from bcing inferred’ and thus the reported limits are assumed to be associated with Sgr A*’s quiescent state. The W. M. Keck Observatory dynamical study of stars in the central stellar cluster (Ghez et al., 1998, 2000; Gezari et al., 2002) provides a rich source of high angular resolution 2 pm data between 1995 and 2002. In, these proceedings, we present 2 pm flux density limits from maps that were each composed of data from a single night. The elapsed time of 3-4 hours in each map is approximately the same as the time scale of the observed X-ray flares, making this data set ideally suited for possibly detecting a flare of this type. Given the quantity of these observations, our non-detections establish a robust upper limit for the flared state’s 2 pm emission intensity.
2 Observations High-resolution, near-infrared observations of the Galactic center were conducted from 1995 June to 2002 July using both speckle and adaptive optics (AO) imaging techniques on the Keck 10 m telescopes. The speckle observations were obtained in the K band (A, = 2.2 pm, AX=0.4 pm) using the Keck I facility near-infrared camera (NIRC; Matthews & Soifer, 1994) with external reimaging optics. This resulted in a pixel scale of 0‘.’0203 and a field of view (FOV) of 5’.!12x5!’12 (Matthews et al., 1996). During each night of observations, several thousand short exposures (tezp= 0.137 sec) were taken in sets of -200. A more limited set of data was collected using two different science cameras behind the Keck Ll A 0 system (Wizinowich et al., 2000b). The first A 0 data set was collected in the K’ band (A, = 2.1 pm, AA=0.35 pm) in early 1999 with the near-infrared engineering camera (KCAM; Wizinowich et al., 2000a), which had a pixel scale of 0!’0175, an FOV of 4!’4~4!’4. Each image had an exposure time of 5 s. The slit-viewing camera of NIRSPEC (SCAM; McLean et al., 1998) provided a second set of A 0 images for this study. These images, like the speckle images, were made in the K band and had a pixel scale of W0170, a FOV of 4!’4 x4!’4, and an exposure time of 10 s. USNO 0600-28579500 served as the natural guide star for all of these A 0 observations. Since this guide star is both faint (R = 13.2) and distant from the target ( r ~ 3 0 ” ) , the A 0 performance was non-optimal. All observations prior to 2002, with the exception of 1998 Aug and 1998 Oct, are reported by Hornstein et al. (2002) and Table 1 provides a summary of all new observations.
3 Data Analysis & Results Three basic steps constitute the data analysis process in this program. First, high angular resolution maps are generated from the individual short exposure frames (53.1). Second, all point sources in the FOV are identified and a direct detection of Sgr A* is ruled out (53.2). Third, limits for Sgr A* are derived from the residual maps, in which all identified point sources have been removed ($3.3). 3.1 Construction of Images
Image processing proceeds similarly to that carried out for the dynamical experiment, with one exception. Rather than combining all the data from the duration of an observing run, typically 2-3 nights, we synthesize the data over each night to produce 38 maps, each of which is limited to an elapsed time of 3-4 hours.
’
While several papers have reported the possible detection of a variable near-infrared source coincident with Sgr A* (Herbst, Beckwirh, & Shure, 1993; Close, McCarthy, & Melia, 1995; Genzel et al., 1997), subsequent high resolution observations have identified this emission to be from high proper motion sources (Eckart et al., 1995; Eckart & Genzel, 1997; Ghez et a]., 1998, 2003).
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Table 1 New Near-Infrared Limits on Sagittarius A*"
PSF Characteristics Limit' Limit" Epoch FWHM,,,, FWHMhalo Eh,,,, (UT) Camera Number of Frames (arcsec) (arcsec) (%) (mag) (mJy) 0.05 0.2 3 16.13 5.1 1998 July4 NIRC 2352 1998 Oct 9 NIRC 2064 0.05 0.3 4 15.40 10.0 0.05 0.3 2 15.55 8.7 2002 Apr 23 NIRC 2880 0.05 0.3 2 16.74 2.9 2002 Apr 24 NIRC 6460 2002 May 23 NIRC 1900 0.04 0.3 3 16.24 4.6 0.05 0.2 5 16.24 4.6 2002 May 24 NIRC 2280 5 16.47 3.7 2002 May 28 NIRC 2660 0.06 0.3 0.05 0.3 3 16.12 5.1 2002 May 29 NIRC 1710 2002 Jun 01 NIRC 5302 0.05 0.2 4 17.09 2.1 0.05 0.3 2 15.61 8.2 2002 J u l l 9 NIRC 1900 2002 Jul20 NIRC 2470 0.06 0.4 3 15.72 7.5 "Used in combination with the results reported previously by Hornstein et al. (2002). bPercentage of total energy contained in the PSF core "Using an A,=30 with Ak/A,=O. 112 and Ak,/A,=O. 1 17 (Melia & Falcke, 2001; Rieke & Lebofsky, 1985), the seventh and eighth columns list observed limits and dereddened values, respectively. Since the details of this method are described elsewhere (Ghez et al., 1998), only a brief summary is provided here. Standard image reduction techniques are applied to all the individual speckle and A 0 frames. For the speckle data, a two stage shift-and-add (SAA; Christou, 1991; Ghez et al., 1998) analysis then produces the final high resolution maps. In the first stage, the 200 frames in each set are combined to form an intermediate SAA image. Then, these multiple intermediate SAA images (from throughout the night) are combined to form one final SAA map. This allows each intermediate image to be examined for seeing quality. In combining the intermediate images, a seeing cut is established so as to exclude those images with the worst seeing from the final map. For the A 0 data, this cut is also carried out on the individual A 0 images before they are registered and averaged together. Figure 1 displays representative final nightly speckle and A 0 maps. 3.2
Point Source Identification & Search for Sgr A*'s near-infrared emission
In all maps, stars are identified using StarFinder, an IDL package developed for astrometry and photometry in crowded stellar fields (Diolaiti et al., 2000). This package iteratively generates estimates of the point spread function (PSF) from a few selected bright stars and then identifies point sources over the entire FOV through cross-correlation of the map with the PSF model. For the PSF extraction, we found that the most reliable PSF models are obtained with a support size of -2", which represents a compromise between needing to accommodate the large PSF halos and yet having a limited FOV. This choice limits our analysis to images with PSF halo sizes of 0!'4 or less, as the PSFs of the remaining two images are poorly characterized by this process. The PSF model is based on four of the five brightest stars in the FOV (IRS 16NE, 16C, 16NW, and 29N; see Figure 1); IRS 16SW is avoided as a PSF model star since it is surrounded by relatively bright stars in its immediate vicinity. For point source identification, a correlation coefficient greater than 0.8 between the PSF model and the actual stellar image is required to avoid spurious detections. This process results in the identification of -100 point sources in each map. The speckle and A 0 images have significantly different PSFs. Nonetheless, both PSFs are composed of a compact core on top of a broader halo. Table 1 provides the characterization of the PSF in each new map based on the radial profile of the PSF model. While the speckle images have PSF core FWHM that are nearly diffraction-limited (-0!'05) and -40% smaller than that of the A 0 images (-U.'08), the A 0 PSF cores contain -30% of the total energy, 12 times more than the typical speckle PSF.
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Fig. 1 Speckle (left) and A 0 (right) images from 1999 May of the Galactic center. The small box indicates a 1”x 1” region centered on Sgr A*, whose approximate position is marked with a cross. Both images are displayed with a histogram equalization stretch to show the fainter stars in the field and are oriented such that north is up and east is to the left.
At this stage in the analysis, it is possible to look for direct detections of Sgr A*. We use proper motion acceleration vectors (Ghez et al., 2000) to pinpoint the central black hole’s position relative to the nominal radio position of Sgr A* (Menten et al., 1997) to within lY.’04 (1 a). Within 0!’08 of this location, 4 sources (13.9 < K < 16.5) are identified, all of which were previously detected in ”monthly” maps made from all data in a single observing run (Ghez et al., 1998,2003) and, furthermore, have significant proper motions. This high stellar density emphasizes the need for improved accuracy in Sgr A*’s position in the IR reference frame in order to measure or constrain its emission. With no stationary source identified in this region, we conclude that Sgr A* has not been detected. 3.3 Flux Density Limits for Sgr A* In order to determine an accurate detection limit at the position of Sgr A*, it is necessary to remove the contaminating seeing halos from nearby sources. A “stars-only” map is created using the PSF model and list of stars generated by StarFinder. This is then subtracted from the original map, producing a residual map. With the residual map, a 3 G point source detection limit for Sgr A* is established based on 3 times the rms of 25 aperture photometry values, which are calculated using -60 mas radius apertures and sky annuli extending from -60 to -90 mas. The 5 x 5 grid of apertures in the residual map corresponds to an area of -0“6x0?6, more than 2 orders of magnitude larger than the uncertainty in the location of Sgr A*. Zero points are obtained through carrying out the same aperture photometry in the original maps (prior to the stars-only subtraction) on all known nonvariable sources brighter than K=lO.S,using the flux densities reported by Blum, Sellgren, & Depoy (1996), and that occur in more than 30% of the frames for each night. The photometric calibration sources used are IRS 16NW, 16C, 16CC, and 16NE, when the FOV allows its inclusion; IRS 29N is omitted since it is found to be marginally variable at the 2 0 level. Typical photometric zero-point 1 a uncertainties of -0.04 mag result from this procedure, Table 1 and Figure 2 contain the resulting 3 a point source detection limits for Sgr A*. The lowest of these upper limits gives an observed limit for the quiescent state of Sgr A* of 0.09~t0.005mJy or, equivalently, a dereddened value of 2.0f0.1 mJy, while the highest limit constrains our analysis of the variable state of Sgr A* to 50.8 mJy (observed) or 19 m l y (dereddened). Although these upper limits
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Fig. 2 The 3 c limiting flux density calculated for Sgr A* for each epoch of observation, including the results reported ) lowest limit of 2.0 mJy and previously by Hornstein et al. (2002) (corrected for reddening by a factor of ~ 2 2 . The the highest limit of 19 mJy constrain Sgr A*'s quiescent and variable states, respectively.
vary significantly from epoch to epoch (because of variable observing conditions and/or a variation in the number of contaminating sources detected and removed), the lowest of them is lower than the best previously reported limit at 2 p m for Sgr A*'s quiescent state prior to Hornstein et al. (2002) (dereddened 4 mJy; Genzel & Eckart, 1999)by a factor of 2.
4 Discussion Using the coverage of the X-ray and near-infraredexperiinents and assuming random 3 hr flaring events that produce both enhanced X-ray and near-infrared emission, we consider the likelihood that a near-infrared flare in excess of our weakest limit occurred over the course of the IR experiment. After an intense X-ray monitoring campaign, the mean rate of factor-of-ten X-ray flares has been established as 0.63~0.3per day (Baganoff et al., 2003). Unfortunately, the flare observed by XMM-Newtonoccurred at the very end of the observations and was not observed in its entirety. Therefore, we do not consider that flare here. Given the duty cycle for the Chandru flares, there was a 10% probability of seeing a flare in any one of our near-infrared observations and a 5% probability of seeing zero flares throughout the entire experiment. We therefore assume that at the 2 (T confidence level, a flare occurred during our experiment, and we use our limits to constrain the variability models. While only a limited amount of modeling of the original 3 hr X-ray flare detected at Sgr A* has been carried out, existing models predict 2 pm dereddened emission as high as 300 mJy in the model preferred by Markoff et al. (2001) but as small as 0.4 mJy by Liu & Melia (2002). The former model explains the elevated X-ray emission, produced by synchrotron self-Compton emission, by an enhanced temperature for the relativistic electron population, as might arise in a magnetic reconnection event. On the other hand, Liu & Melia (2002) present a flare model in which the flare arises as the result of an enhanced accretion rate, and bremsstrahlung emission is dominant at both near-infrared and X-ray wavelengths. Even less modeling has been done of the more recent factor-of-ten X-ray flares (see, e.g., Markoff et al., 2003). For these proceedings, we assume that any factor-of-ten flare produces a detectable near-infrared emission. The lack of a near-infrared detection of Sgr A* makes the Markoff et al. (200 I ) preferred model and any other mechanism that produces flared 2 pm emission in excess of 19 mJy (dereddened) unlikely.
5 Conclusions These proceedings summarize a search for a near-infrared counterpart to Sgr A* in the flared state. From the length of our observations, this search was sensitive to variability on time scales of 3-4 hr. No
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such counterpart was detected. However, by identifying and removing all the stars in the crowded inner -0!’6xU.’6 of the Galactic center, an upper limit for the near-infrared emission from Sgr A* has been inferred for each observation epoch. These limits constrain the quiescent emission from Sgr A* to 50.09 mJy (2.0 mJy, dereddened) and the variable component to 50.8 mJy (19 mJy, dereddened) at the 2 u confidence level. Continued monitoring of Sgr A* at other wavelengths (particularly mid-IR and submm) will provide valuable information about Sgr A*’s flaring properties.
Acknowledgements Data presented herein were obtained at the W.M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California and the National Aeronautics and Space Administration. The Observatory was made possible by the generous financial support of the W.M. Keck Foundation. This work was supported by the NSF through both grant 9988397 and the Science and Technology Center for Adaptive Optics, managed by the University of California at Santa Cruz under cooperative agreement No. AST - 9876783. AMG also thanks the Packard Foundation for its support.
References Baganoff, F. K. et al. 2001, Nature, 413,45 Baganoff, F. K. et al. 2003, these proceedings Blum, R. D., Sellgren, K., & Depoy, D. L. 1996, ApJ, 470,864 Christou, J. C. 1991, PASP, 103, 1040 Close, L. M., McCarthy, D. W., & Melia, F. 1995, ApJ, 439,682 Diolaiti, E., Bendinelli, 0..Bonaccini, D., Close, L., Cume, D., & Parmeggiani, G. 2000, A&AS, 147, 335 Eckart, A. & Genzel, R. 1997, MNRAS, 284,576 Eckart, A., Genzel, R., Hofmann, R., Sams, B. J., & Tacconi-Garman, L. E. 1995, ApJ, 445, L23 Genzel, R. & Eckart, A. 1999, ASP Conf. Ser. 186: The Central Parsecs of the Galaxy, 3 Genzel, R., Eckart, A,, Ott, T., & Eisenhauer, F. 1997, MNRAS, 291, 219 Genzel, R., Pichon, C., Eckart, A,, Gerhard, 0. E., & Ott, T. 2000, MNRAS, 317,348 Gezari, S., Ghez, A. M., Becklin, E. E., Larkin, J., McLean, I. S., Moms, M. (2002), ApJ, 576, 790 Goldwurm, A,, Brion, E., Goldoni, P., Fernando, P., Daigne, F., Decourchelle, A,, Warwick, R. S., & Predehl, P. 2003, ApJ, 584,751 Ghez, A. M., Salim, S., Hornstein, S. D., Tanner, A,, Moms, M., Becklin, E. E., Duchene, G. 2003, ApJ, submitted (asuo-ph/0306130) Ghez, A. M., Klein, B. L., Morris, M., & Becklin, E. E. 1998, ApJ, 509,678 Ghez, A. M., Moms, M., Becklin, E. E., Tanner, A., & Kremenek, T. 2000, Nature, 407, 349 Herbst, T. M., Beckwith, S. V. W., & Shure, M. 1993, ApJ, 411, L21 Hornstein, S. D., Ghez, A. M., Tanner, A., Moms, M., Becklin, E. E., & Wizinowich, P.2002, ApJ, 577, L9 Liu, S . & Melia, F. 2002, ApJ, 566, L77 Markoff, S., Falcke, H., Yuan, E, & Biermann, P. L. 2001, A&A, 379, L13 Markoff, S. et al. 2003, these proceedings Matthews, K., Ghez, A. M., Weinberger, A. J., & Neugebauer, G. 1996, PASP, 108,615 Matthews, K. & Soifer, B.T. 1994, Infrared Astronomy with Arrays: The Next Generation, I. McLean ed. (Dordrecht: Kluwer Academic Publishers), 239 McLean, I. S. et al. 1998, Roc. SPIE, 3354,566 Melia, F. & Falcke, H. 2001, ARA&A, 39,309 Menten, K. M., Reid, M. J., Eckart, A,, & Genzel, R. 1997, ApJ, 475, L1 I 1 Moms, M., Tanner, A. M., Ghez, A. M., Becklin, E. E., Cotera, A., Werner, M. W., & Ressler, M. E. 2001, American Astronomical Society Meeting, 198,4101 Narayan, R. 2002, Lighthouses of the Universe: The Most Luminous Celestial Objects and Their Use for Cosmology Proceedings of the MPAESO,p. 405,405 (astro-pW0201260) Rieke, G. H. & Lebofsky, M. J. 1985, ApJ, 288,618 Stolovy, S. R., McCarthy, D. W., Melia, F., Rieke, G., Rieke, M. J., & Yusef-Zadeh, F. 1999, ASP Conf. Ser. 186 The Central Parsecs of the Galaxy, 39 Wizinowich, P. L.,Acton, D. S., Lai, O., Garthright, J., Lupton, W., Stomski, P.J., 2000a, Proc. SPIE, 4007,2 Wizinowich, P. L. et al. 2000b, PASP, 112, 315 Zhao, J., Bower, G. C., & Goss, W. M. 2001, ApJ, 547, L29
Astron. Nachr./AN 324, No. S1, 377-382 (2003)/ DO1 10.1002/asna.200385041
A New X-Ray Flare from the Galactic Nucleus Detected with XMM-Newton A. Goldwurm*l, E. Brion2, P. Goldoni', P. Ferrandol, F. Daigne', A. Decourchelle', R. S. Warwick3,and P. Predeh14
' Service d' Astrophysique, DAPNIAIDSMICEA, CE-Saclay, F-91191 Gif-Sur-Yvette,France
* Centre d'Etude NuclCaire de Bordeaux-Gradignan,AllCe du Haut Vigneau, 33175 Gradignan, France Department of Physics and Astronomy, University of Leicester, Leicester LEI 7RH, UK Max-Planck Institut fur Extraterrestrische Physik, Postfach 1312, 85741 Garching, Germany
Key words Accretion, accretion disks - Black hole physics - Galaxy: center - X-rays: general PACS 04A25
The compact radio source Sgr A*, believed to be the counterpart of the massive black hole at the Galactic nucleus, was observed to undergo rapid and intense flaring activity in X-rays with Chandra in October 2000. We report here the detection with XMM-Newton EPIC cameras of the early phase of a similar X-ray flare from this source, which occurred on 2001 September 4. The source 2-10 keV luminosity increased by a factor z 20 to reach a level of 4 erg spl in a time interval of about 900 s, just before the end of the observation. The data indicate that the source spectrum was hard during the flare and can be described by simple power law of slope % 0.7. This XMM-Newton observation confirms the results obtained by Chandra, suggests that, in Sgr A*, rapid and intense X-ray flaring is not a rare event and therefore sets some constraints on the emission mechanism models proposed for this source.
1 Introduction The bright, compact radio source Sgr A" is believed to be the radiative counterpart of the 2.6 lo6 Ma black hole which governs the dynamics of the central pc of our Galaxy (Melia & Falcke 2001). The compelling evidence for such a massive black hole at the Galactic Center (see Schodel et al. 2002 for the most recent results), contrasts remarkably with the weak high-energy activity of this object. In spite of the fact that stellar winds and hot gas probably provide enough material for a moderateflow level of accretion, the total bolometric luminosity of the source amounts to less than lop6 of the estimated accretion power (Melia & Falcke 2001, Goldwurm 2001). This motivated the development of several black hole accretion flow models with low radiative efficiency, some of which have also been applied to binary systems, low luminosity nuclei of external galaxies and low luminosity active galactic nuclei. These models include spherical Bondi accretion in conditions of magnetic field sub-equipartition with a very small Keplerian disk located within the inner 50 Schwarzschild radii (Rs),large hot two-temperature accretion disks dominated by advection (ADAF) or non-thermal emission from the base of a jet of relativistic electrons and pairs, and some other variants or combination of the above models. However any such model still predicts some level of X-ray emission from Sgr A* and determining the properties of such emission would constrain the theories of accretion and outflows in the massive black holes and in general in compact objects. The 20 years search for high energy emission from Sgr A* (Watson et al. 198 1, Predehl & Triimper 1994, Goldwurm et al. 1994) has recently come to a turning point with the remarkable observations made * Corresponding author: e-mail:
[email protected],Phone: +3301 6908 2792, Fax: +33 01 6908 6577 @ 2001 WILEY-VCH Verlag GmbH & Ca. KGaA. Weinhem
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with the ChandraX-ray Observatory in 1999 and in 2000. Baganoff et al. (2001a) first reported the detailed 0.5" resolution images obtained with Chandra in the range 0.5-7 keV, which allowed the detection of weak X-ray emission from the radio source. The derived luminosity in the 2-10 keV band was 2 erg scl, for a distance of 8 kpc with steep power law spectrum (index of 2.7) and some evidence that the source is in part extended on a I" scale. Then, in October 2000, the same source was seen to flare up by a factor of = 45 in a few hours (Baganoff et al. 2001 b). The luminosity increased from the quiescent level measured in 1999 to a value of erg s-'. The flare lasted a total of 10 ks but the shortest variation took place in about 600 s, implying activity on length scales of M 20 Rs, for the above quoted mass of the galactic center black hole. Evidence of spectral hardening during the flare was also reported by the authors who determined a source power law photon index during the event of 1.3 (k0.55). significantly flatter than observed during the quiescent state. These results constrain models of the accretion flow and radiation mechanism for Sgr A*. XMM-Newton, the other large X-ray observatory presently in operation, features three large area Xray telescopes coupled to three CCD photon imaging cameras (EPIC) operating in the 0.1-15 keV range and to two reflection grating spectrometers (RGS) working in the 0.1-2.5 keV band (Jansen et al. 2001). Although its angular resolution (6" FWHM) is insufficient for properly resolving Sgr A* in quiescence, an intense flare such as the one seen by Chandra can be easily detected and studied with XMM-Newton. We report here (see also Goldwurm et al. 2003) the detection of such en event during a 26 ks XMM-Newton observation of the Galactic nucleus performed on 2001 September 4 as part of a large Galactic Center survey program with XMM-Newton (Warwick et al. 2003).
2 Results The EPIC data reduction and analysis of this XMM-Newton observation are described in detail in Goldwurm et al. (2003). The image recorded in the central CCD (1 1' x 1 1' for the MOS) was dominated by the diffuse emission of the Sgr A East region, and in order to search for a variable central source we extracted and analyzed light curves from events collected within a 10" radius region centered on Sgr A*. As shown in Fig. 1, the 2-10 keV count rate from the combined MOS 1 and MOS 2 events selected in this way, is quite stable around an average value of 0.08 cts scl till the last 900 s of the observation. Then the count rate gradually increases to reach a value of about a factor 3 higher in the last bin. The integrated count rate in the last 900 s reaches 7 over the average value measured before the flare and the detected variation has a negligible probability to be a statistical fluctuation. A similar peak (4.3 (T)is observed in the counts extracted from the PN camera which stopped observing about 250 s before the MOS (see Fig. 1 b). Similar light curves, from a larger region far from the source do not show any evidence of such an increase in the counting rate. In Fig. 2 we report a MOS image of the region around the nucleus integrated during the 1000 s before the flare and a similar image integrated during the last 1000 s and fully including the source flare. The brightening we detected in the light curves is clearly due to the brightening of a central source. We compared the data to the instrument point spread function to determine the location of the excess. On the 2-10 keV MOS 1 and MOS 2 image of the last 1000 s, rebinned to have pixel size of 4", we obtained the centroid of the source at R.A. (52000) = 17h 4.Y 39.99' Dec (J2000) = -29" 00' 26.7", with a total error, dominated by residual systematic uncertainties in the XMM-Newton focal plane, of 1S".The derived flare position is therefore compatible with the Sgr A* radio location (Yusef-Zadeh et al. 1999), since it is offset from the latter by only 1.5" and we conclude that the flare detected by XMM-Newton is associated with the galactic nucleus. A first, spectral analysis of the flaring event was performed by computing a simple hardness ratio, defined as the ratio between the measured counts (including background) in the hard band 4.5- 10 keV, and those in the soft band 2-4.5 keV. We found that the hardness ratio increased by 0.32 & 0.13 during the event with respect to the value before the flare. Though the hardening has a modest statistical significance
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of 2.5 o,it is fairly consistent with the flare trend observed with Chandra. Since the Chandra data showed that the quiescent emission within 10" from Sgr A"s position contains a dominant diffuse emission and a large contribution from close point sources, we had to model in some way these components to study the flare spectrum (see Goldwurm et al. 2003). We extracted MOS and PN count spectra from the 10" radius circular region centered on Sgr A* before the flare and during the flare (last 900 s for MOS and last 700 s for PN). To derive the spectra reported in Fig. 3 we used the spectra extracted before the flare as background components for the flare spectra. After subtraction of the non-flaring count spectrum, the flaring MOS and PN spectra were fitted with a simple absorbed power law with NH fixed to the Chandra measured value of 9.8 10" cm-2 and leaving the MOS and PN normalizations free to vary. Results both without and including dust scattering are reported in Table I . We obtained a best fit photon index of 0.7 f 0.5 (error at 1 D for one interesting parameter) with x; = 0.98 for 20 d.o.f., that is significantly harder than the spectrum measured with Chandra during the quiescent state (2.7 f l.O), and compatible, within uncertainties, to the index measured during the 2000 October flare. This procedure subtracts from the flare spectrum the non flaring component of Sgr A* and therefore assumes that the quiescent emission from Sgr A* is negligible, This is an acceptable approximation since, if at the level measured in 1999 by Chandra, the quiescent emission contributes by only M 5% to the counts of the flare spectrum. On the other hand, this procedure allows to subtract the diffuse emission present in the region of the spectral extraction and the instrumental background in a model-independent way. We remark that the count excess around 6-7 keV in the residuals of the MOS spectrum of Fig. 3 is not significant. Including a narrow gaussian line at 6.4 keV (with fixed centroid and zero width) to the model of the absorbed power law, we can set an upper limit (90 % confidence level in 1 parameter) to an iron emission line of about 1.8 keV equivalent width. The measured absorbed source flux in the 2- 10 keV band, corrected for the fraction of encircled energy at a distance of 10" (60%),is then of (3.3 f 0.6) 10-l' ergs cmp2 spl (1 F errors by fixing best fit erg s-'. This parameters but normalization), equivalent to a 2-10 keV luminosity at 8 kpc of M 4 is the average value in the last 900 s but the last light curve 180 s bin was about a factor 1.4 higher, thus erg s-l. These numbers are subject to large errors due to the the luminosity reached a value of M 6 low statistics available. But the general result which emerges is that the flare we detected presents a harder spectrum than the one measured with Chandra for Sgr A* during the quiescent period.
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Fig. 1 a) Count rate, sampled in bins of 180 s, collected with both MOS cameras from a region within 10" from Sgr A* in the range 2-10 keV (black upper curve). An equivalent light curve collected from a 30" radius region centered about I' East of Sgr A* and rescaled by a factor 0.1 for clarity, is shown for comparison (red lower curve). Dashed lines indicate the average value computed before the flare. b) Zoom of the Sgr A' MOS light curve (black circles) around the period of the flare compared to a similar light curve (count rate within 10" from Sgr A* in the 2-10 keV band in bins of 180 s) from PN data (red crosses)
A. Goldwurm et al.: Sgr A * X-Ray flare seen with XMM-Newton
380
Fig. 2 Images of the 5' x 5' region around the Galactic nucleus in the band 2-10 keV obtained from MOS events integrated in the lo00 s before the flare (a) and in the last 1000 s of the observation including the flare (b).Pixels were rebinned to a size of 5.5'' x 5.5". Sgr A*'s position is right in the middle of the central bright pixel visible in the flare image (b).
3 Discussion The XMM-Newton discovery of a new X-ray flare of Sgr A* in September 2001 confirms the results obtained in the earlier Chandra observations. XMM-Newton observed only the first part of the flare, but the recorded event is fully compatible in intensity, spectrum and time scales with the early phase of the flare seen by Chandra. This detection of another such large X-ray flare from Sgr A* indicates that the event is not rare. In spite of the lack of detection of flares from Sgr A* in another 50 ks XMM-Newton observation performed in February 2002 (Predehl et al. 2003), the daily presence of X-ray flares in Sgr A* has been recently confirmed by further Chandra observations performed in 2002 (Baganoff et al. 2003). The radio source on the other hand has been observed many times and the detected flux variability has never exceeded a factor 2 (Zhao et al. 2001). This implies that it is unlikely that radio emission presents a comparable large increase in flux and this provides some constraints on the models. The X-ray flare from Sgr A* cannot be explained by pure Bondi or ADAF models (Narayan et al. 1998) as in these models the emission is due to thermal bremsstrahlung from the whole accretion flow and arises from an extended region (between lo3 - lo5 Rs) which cannot account for such rapid variability. Models which predict emission from the innermost regions near the black hole involve a mechanism acting either at the base of a jet of relativistic particles (Markoff et al. 2001) or in the hot Keplerian flow present within the circularization radius of a spherical flow (Melia et al. 2001, Liu & Melia 2002). In both cases a magnetic field is present in the flow and the linearly polarized sub-mm radiation is attributed to optically thin synchrotron emission from the inner region, while the X-rays during quiescent period are produced by the synchrotron self-Compton (SSC) mechanism whereby radio to mm photons are boosted to X-ray energies by the same relativistic or subrelativistic electrons that are producing the synchrotron radiation. However large X-ray flux variations produced by a change in accretion rate in these models would imply an equivalent increase in the radio and sub-mm part of the spectrum not compatible with the lower amplitude of radio changes compared to X-rays (Markoff et al. 2001). Not to mention that the X-ray spectrum would remain rather steep. The model of a inner circularized flow however predicts low or anti correlation of the radio emission with the X-rays if the radiation mechanism for the X-ray flare is bremsstrahlung rather than SSC. The sub-mm and far IR domain on the other hand would in this
Astron. Nachr./AN 324.No. S 1 (2003)
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case show a large correlated increase, but at these frequencies the measurements have not been frequent enough to settle the issue. Though the exact modelling of radiation process depends on viscosity behavior and other uncertain details, the observed hardening of the spectrum during the flare indeed favours the bremsstrahlung emission mechanism in this model rather than the SSC one (Liu & Melia 2002). Another totally different model for the X-ray flares (Nayakshin & Sunyaev 2003) proposes that those result from the interaction of close orbiting stars with a very cold neutral accretion disk around Sgr A*. More compelling constraints on the models will be set when simultaneous observations in radiokub-mm and X-ray wavelengths of such a flare are obtained. We compared the time of the flare to a recent radio light curve of Sgr A* obtained at 1.3 cm and 2 cm with the VLA between 2001 March and November (Yuan & Zhao 2003). The X-Ray flare occurred 1-2 days after a local maximum of the curve, but no radio data points are reported for the day when our XMM-Newton observation took place. It will b e also important to study the shape of the flare spectrum at energies higher than 10 keV to fully understand the radiation mechanism producing the X-rays. In particular by measuring the high energy cut-off of the spectrum one could determine the electron temperature for a thermal emission or the Lorentz factor for non-thermal processes. We estimated that such a flare should be marginally visible in the range 10-60keV with the low energy instruments on board the new gamma-ray mission INTEGRAL,, operating since 2002 October, if the spectrum extends to these energies with the slope observed with Chandra and XMM-Newton. Acknowledgements This paper is based on observations with XMM-Newton, an ESA science mission with instruments and contributions funded by ESA member states and the USA (NASA). References Baganoff, F., et al., 2001a. ApJ, in press (astro-pW0102151) Baganoff, F., et al., 2001b, Nature, 413, 45 Baganoff, F., et al., 2003, these proceedings Goldwurm, A., et al., 1994, Nature, 371,589 Goldwurm, A., 2001, Proc. of the 4th INTEGRAL Workshop, ESA-SP 459,455 Goldwurm, A., et al., 2003, ApJ, 584,751 Jansen, F., et al., 2001, A&A, 365, L1 Liu, S., & Melia, F., 2002, ApJ, 566, L77 Markoff, S.,et al., 2001, A&A, 379, L13 Melia, F. & Falcke, H., 2001, ARAA, 39, 309 Melia, F., Liu, S., Coker, R. F., 2001, ApJ, 553, 146 Narayan, R., et al., 1998, ApJ, 492, 554 Nayakshin, S. & Sunyaev, R., 2003, MNRAS, submitted (astro-pW0302084) Predehl, P. & Triimper, J., 1994, A&A, 290, L29 Predehl, P., et al., 2003, these proceedings Schodel, R., et al., 2002, Nature, 419, 694 Watson, M.G., et al., 1981, ApJ, 2.50, 142 Warwick, R. S., et al., 2003, these proceedings Yuan, F. & Zhao, J., 2003, Chin. J. Astron. Astrophys., in press (astro-pW0203050) Yusef-Zadeh, F., Choate, D., Cotton, W., 1999, ApJ, 51 8, L33 ZhdO, J., Bower, G. C., GOSS, W. M., 2001, ApJ, 547, L29
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Table 1 Spectral Fit to X-ray Emission from within 10" from Sgr A* during the Flare
Power-law Model
NH [loz2 cm-']
No Dust Scattering
Dust Scattering
9.8
5.3
Photon Index
0.7
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Norm MOS [10W4 ph cmP2 sP1keV-']
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Norm PN [lop4 ph cm-' s-' keV-']
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0.2;;;:
x: (d.0.f.) 0.98 (20) Notes : Scattering computed for fixed value of Av = 30. Normalization is the flux density at 1 keV. Errors are at 68.3% confidence interval for 1 interesting parameter.
I
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1
I
I
2
0.95 (20)
5
I
10
Energy (keV) Fig. 3 Count spectra from MOS (black data point set) and PN (red data point set) data, extracted from a region of 10" radius around Sgr A* during the flare after subtraction of the non flaring spectra. The spectra are compared to the best fit model of an absorbed power law without dust scattering (see Table I )
Astron. Nachr./AN 324, No. S1,383 -389 (2003) I DO1 10.1002/asna.200385090
Searching for Structural Variability in Sgr A* Zhi-Qiang Shen*',5,6,M. C. Liang'.*, K. Y. Lo3, and M. Miyoshi4
'
Academia Sinica Institute of Astronomy and Astrophysics, PO Box 23-141, Taipei 106 Division of Geological and Planetary Sciences, California Institute of Technology, Pasadena, CA 91 125 National Radio Astronomy Observatory, 520 Edgemont Road, Charlottesville, VA 22903 National Astronomical Observatory Japan, Osawa 2-21-1, Mitaka, Tokyo 181-8588 Shanghai Astronomical Observatory, Chinese Academy of Sciences, Shanghai 200030 National Astronomical Observatories, Chinese Academy of Sciences, Beijing 100012
Key words Galaxy: center, galaxy: individual (Sagittarius A*), techniques: interferometric Abstract. A model fitting procedure for estimating the parameters of an elliptical Gaussian model that describes the radio emission of Sgr A' observed by the VLBA is presented. By (implicitly) using the amplitude closure relation while fitting the amplitude, the procedure can minimize the calibration errors in millimeter wavelength VLBI measurements, which are crucial to our ongoing effort to search for the structural variability in Sgr A*. The preliminary results from the application to seven-epoch A7 mm VLBA observations seem to show the sign of change to the source apparent structure in at least one epoch over the period of 7 years (1994-2001). Because of the large uncertainties in the determination of minor axis caused mainly by the poor spatial resolution along the north-south direction with the VLBA, these results are suggestive but not conclusive. This demonstrates the necessity of adding the NRAO GBT antenna to the future A7 mm VLBA observations, which can greatly improve the resolution in the north-south by a factor O f 3.
1 Introduction Sgr A*, the extremely compact radio source at the Galactic Center, is the best candidate for a single massive black hole from both the observational results and the theoretical models. Recent significant progress on the observations of stellar dynamics of the Galaxy's central stellar cluster has provided new compelling evidence for the existence of a compact dark mass of 3.0 x 10' Mawithin the vicinity of Sgr A*(Schodel et al. 2003; Ghez et al. 2003). Improvement on the determination of the upper limit to the absolute proper lo5 Mo(Reid et motion of Sgr A* has also placed a stringent constraint on its mass t o be greater than al. 2003). The total flux density variation has been puzzling ever since the discovery of Sgr A* almost three decades ago (Balick & Brown 1974). In 1982, Sgr A* was first reported to be variable at A1 1 and 3.7cm radio wavelengths (Brown & Lo 1982). Since then, there have been a lot of intensive monitoring observations using all the available radio interferometers, such as the Very Large Array (VLA) at A20 to 1.3 c m and 7 mm, the Green Bank Interferometer (GBI) at A1 1 and 3.6 cm, the Nobeyama Millimeter Array (NMA) at A3 and 2 m m and the Sub-Millimeter Array (SMA) at X1.3 and 0.87mm. As a result, the total flux density variation has been seen at all the observing wavelengths from centimeter to submillimeter on all time scales from years to days. It was also found that the variation appears to peak at shorter wavelength first with a relatively larger amplitude fluctuation. Recent analysis of radio light curves seems to suggest a quasi-periodic oscillation, or a double quasi-periodic oscillation (Zhao et al. 2003). N
N
* Corresponding author: e-mail:
[email protected].!w, Phone: +886-2-3365-2200, Fax:+886-2-2367-7849
@ 2003 WILEY-VCH Verlag GmbH & Co KGaA, Weinhem
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3 84
The discovery of a very strong, highly variable and short-lived X-ray flare from Sgr A* by the Chan~ Ma dark masses into a region less than 20 times its dra X-ray Observatory (CXO) places the 3 . 0 lo6 Schwarzschild radius (Baganoff et al. 2001). Intriguingly, 10 days after the X-ray flare, a relatively low amplitude variation was detected by the VLA (Zhao 2002). Nowadays, there is no doubt that Sgr A* is a temporally variable source. The observed intensity variations at both the radio and X-ray bands might be correlated and both might arise from the instabilities in an accretion disk. Thus, it would seem likely that the variability in the flux density of Sgr A* would be accompanied by structural changes in Sgr A*, i.e., Sgr A* should be variable spatially as well.
1 2h
1qh
16'
1gh
20h
- VLBLPT Fig. 1 Uptime plot showing the elevation angle of Sgr A* at different time (in local sidereal time (LST) at PT) seen at the 7 VLBA antenna sites (labelled),from which the fringes are consistently detected. It is clear that the VLBA always looks at Sgr A* at a low elevation, which severely limits the resolution along the north-south direction. LST
Very Long Baseline Interferometry (VLBI) technique is proved to be the most effective and powerful tool for investigating the high-resolution structure of compact objects (Kellermann & Moran 2002). Attempts to measure Sgr A* structure with the VLBI observations, however, have suffered from the angular broadening caused by the diffractive scattering by the turbulent ionized interstellar medium, which dominates the resultant images with a A2-dependence apparent size (Lo et al. 1997). Over the past decade, VLBI experiments have been carried out steadily at millimeter wavelengths (A7 and 3.5mm). Due to the southerly declination of Sgr A* (- - 30"), and the high northern latitudes for most of the existing millimeter VLBI antennas, much of the observational data were taken at low elevation angles where atmospheric effects are substantial (see Fig. I). This fact imposes two limits on the VLBI study of Sgr A*. One is related to the spatial resolution. It would limit the (u,v) sampling in the north-south direction, which happens to be along the minor axis of the scattering structure. Consequently, the spatial resolution in the north-south direction is always inadequate as compared to the scattering size. The other effect is upon the data calibration. The atmospheric absorption due primarily to spectral line transitions of water vapor and oxygen at millimeter wavelengths increases with decreasing elevation angle (larger opacity at lower elevation). The compromised sensitivity or lower signal to noise ratio (SNR), when combined with the short and variable coherence time at millimeter wavelength, results in large calibration uncertainties. As a result, any possible variation in the observed source structure has so far been ascribed to the errors in the calibration of the VLBI data. In order to minimize the calibration errors and thus to improve the accuracy of the measurements, we have developed a model fitting procedure by using the amplitude closure relation. The preliminary results from the application of this method to the existing A7 mm VLBA observations seem to show the sign of
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change to the source apparent structure in at least one of seven epochs. However, the significance level of such a change is seriously limited by the poor spatial resolution along the north-south direction with the current VLBA. The situation can be improved by adding the NRAO GBT antenna to the future VLBA observations at X7 mm.
2 Description of the algorithm To improve upon the calibration of current VLBI observations of Sgr A*, we adopted and implemented a model fitting procedure in which the amplitude closure relation is applied. By measuring the closure amplitudes, the knowledge of the absolute flux density is lost, as the absolute position in the case of the closure phases.
2.1
Two basic equations
In general, by ignoring the baseline-dependent effects (in complex gain and noise), the relationship between the observed visibility TbS ( t )and the true source visibility qyueon a baseline i j can be expressed as
where
G, ( t ) G3 , ( t ) is the time-variable complex gain of antenna i and j, respectively, qzJ( t ) is a stochastic thermal noise having a zero mean and variance cz". The best-fit model of source structure can be obtained by minimizing x2, the weighted sum of square4 of the residuals qz3(t),between the observed data l$bs(t) and the predicted data G,(t)G,*(t)V,Y'(t), for all available baselines throughout the whole observation, defined as below,
where ui,,~(t) is the weighting coefficient which is chosen to be the inverse of the
variance of the data due to noise, i.e., 'UIQ= -&, gv
yy"d(t)is the visibility from the source structure model representation. These equations are similar to those widely used in the self-calibration technique for VLBI imaging (Pearson & Readhead 1984). Both are, by definition, consistent with the closure quantities (closure phase and closure amplitude). Both procedures first use the same method to determine the complex antenna gains as a function of time. The difference between two methods lies in the way of obtaining the model of the source sky brightness distribution. Both algorithms should converge to the consistent results with the high enough SNR data. The conventional imaging process has been dominated by the CLEAN deconvolution method, whose biggest drawback is non-uniqueness in the final image, especially when the SNR is poor. The model fitting algorithm, however, can easily search over all the possible models. This could be very effective in fitting to a model that has fewer free parameters as is the case of Sgr A*, whose radio emission, as a first order approximation, can be well represented by an elliptical Gaussian model (3 parameters only). This is because most of the closure phases measured by VLBI observations of Sgr A* are consistent with zero (e.g. Doeleman et al. 2001). Furthermore, the fact that the scattering image of X2-dependent size is always resolved out on the short to intermediate, depending on the observing wavelength, baselines of the VLBA means that we deal with a lot of weak detections ( S N R decreases with increasing projected baseline length).
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2.2 Model visibility amplitude From now onwards, our discussion will be restricted to a single Gaussian model whose brightness distribution is symmetric with zero visibility phase. Assuming it is elongated along a position angle (East of North) in degrees with an axial ratio CY and a major axis size of 0 in radians, its normalized visibility amplitude V;jnod also has a Gaussian distribution as
TI)distance of baseline i j in wavelengths, where pij is (u,
paj = 4.2
(ual~
0 - vlJ s szria)2 ~
+ (u,,s z n +~ v,,
c,05+)2
14)
~
here, uzj and ubJ are the East and North components, respectively, of the projected baseline vector seen from the source, in the units of wavelengths.
2.3 The determination of the time-variable antenna-dependent gains As mentioned above, exactly the same self-calibration algorithm is used to solve for the time-variable antenna-dependent complex gains. This usually means an iterative (non-linear) process in order to find gains by minimizing the weighted sum of squares of the residuals v,(t), between the observed data L$‘(t) and the predicted data G,(t)G;(t)Kyod(t), for the available baselines within each integration period instead of the whole observation. For Sgr A* with a symmetrical Gaussian model (zero phase), the antenna-dependent gains should be real and therefore can be expressed as G,(t) = e g z ( t f . By doing so, we can rewrite Eq. ( I ) as a linear function of g2( t )
= +)VV”,’”p‘ ’ ((tt)) N variance ( ~ 2 ’ ~ = ) ~ (+)’v,:. ’ ( t )
where X,,(t)
e, VOb9(t)
and a;,(t)
=
Oa3
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y y ” d ( t ) X“,(t) -
Thus, instead of solving the non-linear least squares problem, we can
easily solve the normal equation (Eq. ( 5 ) )for gi(t) (and gain G,(t))for Sgr A* using the standard matrixinversion method. In matrix notation, regarding time dependence as implicit we have
9 =
( A ~ ~ -lATw’ A ) Y,
where “ + I ” refers to the inverse matrix operator, T to the transpose operator, and g = [gl, 9 2 , ..., g7nlTx,
(m: the number of antennas),
Y = 1nX = [InXlz, InXls, lnXz3,
A is a steering matrix 1 1 0 0 ... 0 0 0 0 1 0 1 0 ... 0 0 0 0
A=
. . . . 0 0 0 0 ... 0 1 0 1 2
x m
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w' is the diagonal matrix of weighting functions as
2.4 Measured visibility amplitude The outputs from the VLBI correlator are complex ones that are integrated over To (typically, 2 seconds for continuum observations of Sgr A*). After global fringe fitting, these are first coherently averaged over T, (10-20 seconds) to obtain a mean visibility amplitude Z k . The baseline denotation i j is now regarded as implicit, for simplicity. Unlike the original complex correlator output, the measured visibility amplitude will have a positive bias with respect to the true visibility amplitude (Thompson, Moran & Swenson 1986). The magnitude of such a bias is a function of SNR with the strongest bias occurring at low SNR. Therefore, in order to use the amplitudes alone to fit a symmetrical Gaussian model to Sgr A* data which have a lot of low SNR ( 5 3 ) detections on long baselines, we must correct for this bias first. Two methods, each involving an incoherent average, were adopted to carry out such a bias-correction, Method one is to, for each observing scan T, (typically, 6-10 minutes) that consists of N( =
Ivy + 2
2;
(7)
< z4> = IV14 + 81V122 + 8cr4, where (T is the rms noise level in 21, and assumed to be constant within T,, and < Z 2 > are the arithmetic mean of 2;and Zt,respectively, < Z 2 >= Z; and < Z4 >= Thus, we can obtain
cf,
J V I = [2 < z2>2 - < 2
4
>ill4 ;
(8) and < Z4 >
& CF='=, 2:. (9)
Method two is to solve for an unbiased estimate of the amplitude directly by doing incoherent average within each scan T, (Rogers, Doeleman & Moran 1995; Doeleman et al. 2001)
r.
N
with the variance in the bias corrected amplitude as
here CQ is the noise in 2,which is a constant (inversely proportional to the square root of the product of T, and the bandwidth) in the unit of the correlation coefficient and, the corresponding SNR value sk is defined as 2,/uI,.
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2.5
The practice of model fitting and error estimation
In practice, only the visibility amplitude information is used for a fitting to the VLBI observations of Sgr A*, i.e., both measured and model visibility vectors are replaced by their amplitudes in solving for the least squares minimization Eq. (2). To find a best fit model that corresponds to the minimal x:~,, a thorough search for x2 in a 3-dimensional space that covers a wide combination of three model parameters (major and minor axes and the position angle) is carried out. In principle, this procedure is applicable to any source whose brightness distribution can be represented by a single Gaussian. should be distributed For the error estimation, we first draw attention to the presumption that &%, as a chi-square distribution with N,l,f degrees of freedom to ensure a credible fit. By definition, this is equivalent to the statement that the best fit model should have reduced chi squares, x ~ = x ~ , , / N d ,equal f, However, in many cases the value for xz at is significantly to unity with a deviation of larger than 1.0 (c.f. Bietenholz et al. 1996). So, we scale up the 68.3% confidence region of parameter + Ax2 with space, as an increase of x2 from xk,, to xLLrL
d m .
xk,,
instead of Ax2= 1. Here, Nd,f is estimated by the summation of the difference between the number of visibilities Nu,, and the number of antennas Nan, over all scans, minus the number of fitting parameters (3 in case of Gaussian model), i.e., N d , f = C,(N,i, - N,,t) - 3. By projecting this confidence contour onto the axis of parameter of interest, we can finally obtain the 1 standard error for that single parameter.
3 Application to the A7 mm VLBA observations of Sgr A* We have applied this model fitting method to seven existing A7 mm VLBA observations of Sgr A* over a period of 7 years (1994 - 2001). The preliminary results show that the apparent source size is dominated by the scattering effect with a position angle (PA) -75". The major axis can be well determined with an error of 0.01 mas, while the minor axis has a quite large error 20.05 mas. Of 7 epochs, however, six epochs showed that the fitted sizes of both major and minor axes are all above the inferred scattering size. On May 31, 1999, a consistent increase in the apparent source size over the scattering size along the minor axis can be seen at all three simultaneously observed A7 mm bands of 39, 43 and 45 GHz with the corresponding significance levels of 3u, 2u and 2a, respectively (see Fig. 2). At another epoch (July 3 1,2001) which was just three weeks after an SMA flare towards Sgr A* (Zhao 2002), the source seems to undergo a structural variation too (Miyoshi et al. 2003), but the uncertainty of 0.12mas is just too large. At this moment, we consider these results to be suggestive.
4 Discussion For two epochs in April and May 1999 when the time-dependent relative gains as a function of elevation could be determined through a least-squares fit of the auto-correlation spectra of the SiO maser towards VX Sgr at different elevation to a well calibrated total power spectrum, we also tried self-calibration imaging and found the consistent results with those from the model fitting. This demonstrates that in principle the difficulties due to the atmosphere can be minimized either by a careful calibration using interleaved maser line observations or by model-fitting using the closure quantities. The resolutions which can be reliably achieved by the VLBA observations of Sgr A* is about I .3 mas by 0.5 mas along PA=10" at A7 mm, compared to the scattering size of 0.37mas by 0.69 mas at PA=8Oo. Obviously, the resolution along the north-south (1.3 mas) was too poor to resolve the scattering size of 0.37 mas. This can explain why the major axis of Sgr A* at a PA-75" can be determined fairly well while the minor axis has a larger uncertainty. Tests on the simulated data also confirm that the lack of
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Fig. 2 The plot of the measured Sgr A* angular size versus the observing frequency. The open circles and filled circles are the sizes of major and minor axes measured at three simultaneously observed A7 mm bands of 39, 43 and 45 GHz by the VLBA on May 31, 1999. The solid line and dashed line represent the scattering dominated sizes as a function of the observing frequency along the major and minor axes with a frequency dependency of 1287/u$~,and 684/uKH, mas (Lo et al. 1998). A consistent deviation in the apparent source size from the minor scattering size at all three bands can be seen.
resolution is an insurmountable barrier to the current VLBA in the determination of minor axis with a good enough accuracy. We have proposed to add the recently commissioned NRAO GBT antenna to the future A7 mm observations. This will improve the resolution along the north-south direction by a factor of 3 by connecting to an isolated, north-south oriented HN-SC baseline. Strong signals have been consistently detected on this single baseline in five of seven epochs data sets (Shen et al. 2003, in preparation). However, the non-detection between either of SC and HN with the rest of the VLBA handicapped the uses of SC-HN detection in the self-calibration process. By introducing the GBT antenna that is available at A7 mm, it is certain that these SC-HN detections can be brought back as the GBT will serve as a “bridge” between SCHN and those western VLBA antennas (e.g., NL/FD).
References Baganoff, F. K., et al. 2001, Nature 413,45 Balick, B, & Brown, R. L. 1974, ApJ 194,265 Bietenholz, M. F., et al. 1996, ApJ 457, 604 Brown, R. L., & Lo, K. Y. 1982, ApJ 253, 108 Doeleman, S. S., et al. 2001, AJ 121, 2610 Ghez, A. M. et al. 2003, these proceedings Kellennann, K. I., & M o m , J. M. 2001, ARAA, 39, 457 Lo, K. Y., et al. 1998, 508, L61 Miyoshi, M., et al. 2003, these proceedings Pearson, T. J., & Readhead, A. C. S. 1984, ARAA, 22,97 Reid, M. J., et al. 2003, these proceedings Rogers, A. E. E., Doelernan, S. S., & Moran, J. M. 1995, AJ 109, 1391 Schodel R., et al. 2002, Nature 419, 694 Thompson, A. R., Moran, J. M., & Swenson, G. W., Jr. 1986, Interferometry and Synthesis in Radio Astronomy. New York: Wiley-Interscience. First (1991) and second (1994) reprintings by Krieger Pub. Co., Malabar (Florida) Zhao, J.-H. 2002, in GCNEWS, Vol. 15,4 Zhao, J.-H., et al. 2003, these proceedings
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Astron. NachrJAN 324, No. S1,391-395 (2003) / DO1 10.1002/asna.200385113
Observations of the Galactic Centre at 610 MHz with the GMRT Subhashis Roy * and A. Pramesh Rao ' National Centre for Radio Astrophysics (TIFR), Pune University Campus, Post Bag No.3, Ganeshkhind, Pune 41 1 007, India. Abstract. We have observed the central region of the Galaxy at 610 MHz with the Giant Metrewave Radio Telescope (GMRT). We detect emission from Sgr A*, the compact object at the dynamical centre of the Galaxy, and estimate its flux density at 610 MHz to be 0.52~0.1Jy. This is the lowest frequency at which Sgr A* has been detected. Comparison of the 610 MHz and 1.4 GHz map of this region indicates that most parts of the Sgr-A West HII region are optically thick at 610 MHz, having optical depth of 2. However, Sgr A*, which is seen in the same region in projection, shows a flat spectral index between 1.4 GHz and 610 MHz. This is consistent with its high frequency spectral index, which indicates that Sgr A* is located in front of the Sgr-A West complex. An intrinsic turnover in its spectrum between 580 and 330 MHz is likely to be the cause of its non detection at 408 and 330 MHz. N
1 Introduction The Galactic Centre (GC) region has been observed at radio wavelengths with high resolution by the Very Large Array (VLA) at 2 cm (Yusef-Zadeh & Wardle 1993) 6 cm (Ekers et a]. 1983), 20 cm (Pedlar et al. 1989) and 90 cm (Pedlar et al. 1989; Anantharamaiah et al. 1991;LaRosa et al. 2000) and several sources were identified within the central half a degree region of the GC. At the dynamical centre of the Galaxy ~ black hole candidate (Ghez et al. 1998), which coincides with a compact nonthermal is the 2 . 6 ~ 1 0Ma radio source named Sgr A*. Around this point source in a somewhat larger scale in projection, are the three arm spiral configuration of ionised gas and dust known as Sgr-A West (Ekers et al. 1983). Near Sgr-A West, is the supernova remnant Sgr-A East. A 7' halo which is thought to be a mixture of thermal and non-thermal emission (Pedlar et al. 1989) can also be identified in this region. Sgr A* has attracted considerable attention from the time of its discovery (Balick & Brown 1974) since it is associated with the nearest supermassive black hole and could be a prototype for such black holes in extra galactic AGNs. This object has now been studied from radio to the X ray ranges (see Melia & Falcke 2001, and references therein). Though the observational data strongly associates it with the 2 . 6 lo6 ~ Ma black hole at the centre (Genzel et al. 1996; Ghez et al. 1998; Reid et al. 1999), there are several questions that remain unanswered. CornpaTed to AGNs, this object is extremely underluminous at all wavelengths, radiating at x lo-'' times of its Eddington luminosity. It is known to vary at higher radio frequencies, and the flux density variations appears to have a periodicity of 106 days (Zhao et al. 2001). Though, no linear polarisation has been detected at radio frequencies, circular polarisation from this object has been detected (Bower et al. 1999). Sgr A* has not been detected below 1 GHz and observations by Davies et al. 1976 at 408 MHz and Pedlar et al. 1989 at 330 MHz provide upper limits on its flux density. It could have a low frequency turnover below 1 GHz, but the nature of the turnover has never been clarified in detail (Melia & Falcke 200 1 ). In order to estimate the spectrum of the Sgr A* at low radio frequencies, we observed it using the Giant Metrewave Radio Telescope (GMRT) at 620 MHz in Aug & Sep 2001 and at 580 MHz in June 2002. In * Subhashis Roy: e-mail:
[email protected] @ 2003 WREY-VCH Verlag GmbH & Co. KGaA. Wemheim
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this band, the free-free optical depth of the Sgr A West is expected to be moderate, and it should be possible to identify Sgr A* in the Sgr A West complex. At these frequencies, the GMRT has an angular resolution of about 6'' and field of view 44'. The details of these observations and the data analysis will be described elsewhere (Roy & Rao 2003, in preparation).
2 Results
Fig. 1 The central 15' region of the Galaxy at 610 MHz as observed by the GMRT. The resolution of the image is 1I .4"x 7.6", with the beam position angle of 7'. The RMS noise is about 6.5 m l y beam-'.
Fig. 2 4.8 GHz continuum map of the Sgr A complex (Yusef-Zadeh 1989) in contour overlaid on the 610 MHz gray scale map of the the same region.
2.1 Features in the 610 MHz map The 610 MHz map of the central 15' region of the Galaxy is shown in Fig. 1. The compact source Sgr A* is clearly seen along with other well known sources like Sgr A West, Sgr A East and the 7' halo. The prominent non-thermal filamentary structure, the Radio Arc is clearly seen in a map covering a larger field. An emission feature =30" south of Sgr A* can be seen in the 610 MHz gray scale map. This feature was identified by Pedlar et al. (1989), who suggested that it is associated with Sgr A East. In order to compare the smaller scale features near Sgr A West with what is seen at higher frequency, we have plotted in Fig. 2 the 4.8 GHz VLA map of this region (Yusef-Zadeh 1989) on the gray scale GMRT image with both maps convolved to the same resolution. Sgr A* is clearly seen in both the maps. There is almost one to one correspondence between the higher emission features at 4.8 GHz comprising the Sgr A West region and the relative lower emission features (indicated by white region in the gray scale map) seen at 610 MHz.
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Fig. density ratio Of the map to the GMRT 6 10 MHz map
GHz VLA GC
Fig. 4 The spectrum of Sgr A* from 300 MHz to 20 GHz. Except the 330 Pedlar et al. (1989) and 610 MHz measurements, all the other data points are taken from Melia & Falcke (2001)
2.2 Optical depth of the ionised gas near Sgr A West In order to study the free-free absorption by the ionised gas, we have divided the 1.4 GHz map (made from the archival VLA data acquired and presented by Pedlar et al. 1989) by the 610 MHz map (both the maps were made with same resolution), which is plotted in Fig. 3. In most parts of Fig. 3, this ratio is < 1 confirming that the emission from the 7' halo and Sgr A East is non thermal and probably due to synchrotron emission. However, in the central region comprising Sgr A West and the 3 HI1 regions near the eastern boundary of Sgr A East the ratio is -2 indicative of free free absorption. In the vicinity of Sgr A*, the radio emission has contributions from a non-thermal component mainly due to Sgr A East and a thermal component due to emission from Sgr A West. Since Sgr A West can absorb the background emission as well as its own emission, separating the two components are essential to estimate the optical depth of the free-free absorption. If the 610 MHz emission towards Sgr A West is dominated by thermal emission, then the optical depth of thermal emission can be estimated from the formula T b = Te.[l- eP7]where T b is the brightness temperature, T, is the electron temperature and 7 is the optical depth. Since the flux density at 1.4 GHz near Sgr A* is a factor of two higher than at 610 MHz, the estimated optical depth in the case of only free-free self absorption is ~ 2To. estimate the non-thermal emission at 610 MHz, we use the GC maps made at 3.6 (Roberts & Goss 1993) and 6 cm (Yusef-Zadeh 1989). Emission from the GC region at wavelengths equal to or less than 6 cm is mostly thermal and is believed to be optically thin. Therefore, based on the 3.6 cm map, we constructed a model for the thermal emission from Sgr A West at 6 cm and subtracted it from the 6 cm map. The difference map provides the excess non-thermal emission at 6 cm as compared to 3.6 cm. With an assumed spectral index of - 1.0 for the non-thermal emission, we estimated the non-thermal component at 610 MHz. After considering this non-thermal component along with the thermal emission from Sgr A West, we estimate an optical depth of 2.6 d10.5 for the ionised gas seen towards Sgr A*.
394
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Estimated flux density of Sgr A*
To estimate the flux density of Sgr A*, we partially resolved out the extended emission around it by applying a shorter uv cutoff of 7 ICX while imaging. The flux density estimated from the image plane is about 0.5 10.1 Jy. We note that in an image of this region even after applying a short uv cutoff, there is significant background confusion in a beam of size 7.5” x 4”.This confusion causes an uncertainty of about 0.1 Jy in the estimated flux density. Therefore, we also estimated the flux densities of this object from the uv plane. The estimated flux density after fitting a elliptical Gaussian model to S g r A* is also 0.5 kO.1 Jy. The respective major and minor axis of the Gaussian fit is 3.8”k0.4” and 1.8”10.6/’ with a position angle of 9 3 f 4 ” . We note that the estimated size of Sgr A* at 610 MHz appears to be consistent with what is expected from scatter broadening (Lo et al. 1998) at this frequency.
3 Discussions 3.1 Low frequency spectral index of Sgr A* While the high radio frequency spectrum of Sgr A* is well established, the spectrum below 1.4 GHz is not well determined. At 1.4 GHz, the flux density of Sgr A* (Zhao et al. 2001) is about 0.5 Jy and its spectral index between 1.4 and 8.5 GHz is +0.17 (Melia & Falcke 2001). Davies et al. (1976) found the flux density of Sgr A* to be a factor of 2 less than at 1.6 GHz and suggested that it has a low frequency turnover around 1 GHz. This appeared to be confirmed from their upper limit of 50 mJy to its flux density at 408 MHz and the 100 mJy upper limit set by Pedlar et al. (1989) at 330 MHz. Our measured flux density of 0.5 Jy at 610 MHz raises questions about the earlier measurements. Measurements of the flux density of Sgr A* between 1.4 GHz and 610 MHz which were close in time show that the spectrum between these frequencies is consistent with that between 1.4 and 8.5 GHz. The spectrum of Sgr A* from 300 MHz to 20 GHz is shown Fig. 4. The 610 MHz observations cover a span of nearly an year and no significant propagation effects due to the Inter Stellar Scintillation (ISS) could be detected. Our measurements rule out any turnover at around 1 GHz and indicate that the turnover has to be at frequencies less than 580 MHz. However, the upper limits at 408 and 325 MHz pose problems for this picture since this would require the spectrum to fall steeper than can be explained by thermal free free absorption or synchrotron self-absorption. We have re-examined the 408 MHz upper limit by Davies et al. (1976) and found that their analysis has not taken account for scatter broadening of Sgr A* (8“ at 408 MHz) and the corrected upper limit at 408 MHz is as high as x 2.5 Jy. The 330 MHz upper limit, however is reliable and implies an inverted spectrum with spectral index > 2.5, suggesting that more than one process could be responsible for the turnover. The nature of the low frequency turnover of Sgr A* could occur due to synchrotron self absorption, internaUexternal free-free absorption or Razin effect. It is of great interest to understand the nature of Sgr A* and its environment, and more careful and systematic flux density measurements below 600 MHz are required. These observations should be spread over many years to eliminate the uncertainty due to interstellar scintillation which could have a large time scale in this direction. 3.2
Location of the Sgr A*
At 610 MHz, Sgr A West shows evidence for free free absorption. In the previous section, we have estimated the optical depth of this ionised gas in the Sgr A West region to be x2.6. If Sgr A* would have been located behind Sgr A West, then its flux density would have been attenuated by a factor of 10. However, the spectral index of Sgr A* between 610 MHz and 1.4 GHz is roughly the same as between 1.4 GHz and 8.5 GHz and shows no effect of the free free absorption by Sgr A West. This indicates that S g r A* is located in front of Sgr A West. It is possible to have alternate scenarios like sharp enhancement of the 610 MHz flux density of Sgr A* to compensate for the absorption due to Sgr A West or a hole in Sgr A West along the line of sight to Sgr A*, but these appear to be unlikely. Sgr A* is located slightly north of the junction of the northern and the eastern arm of the Sgr A West. We have inspected the 3.6 cm map
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(resolution x 0.5”) of the region (Roberts & Goss 1993), and w e do see weak emission likely to be from the diffuse ionised gas which becomes optically thick at 610 MHz. Any hole in this ionised gas has to be smaller than 1” which is unlikely. Thus, Sgr A* is either located in front of Sgr A West, or is embedded within the ionised gas such that the optical depth of that region is 20.1 at 610 MHz.
4 Acknowledgements We thank the staff of the GMRT that made these observations possible. GMRT is run by the National Centre for Radio Astrophysics of the Tata Institute of Fundamental Research.
References Anantharamaiah, K. R., Pedlar, A., Ekers, R. D., & Goss, W. M. 1991, MNRAS, 249,262 Batick, B. & Brown, R. L. 1974, ApJ, 194,265 Bower, G. C., Falcke, H., & Backer, D. C . 1999, ApJL, 523, L29 Davies, R. D., Walsh, D., & Booth, R. S. 1976, MNRAS, 177, 319 Ekers, R. D., van Gorkom, J. H., Schwarz, U. J., & Goss, W. M. 1983, A&A, 122, 143 Genzel, R., Thatte, N., Krabbe, A., Kroker, H., & Tacconi-Garman,L. E. 1996, ApJ, 472, 153 Ghez, A. M., Klein, B. L., Moms, M., & Becklin, E. E. 1998, ApJ, 509, 678 LaRosa, T. N., Kassim, N. E., Lazio, T. J. W., & Hyman, S. D. 2000, AJ, 119, 207 Lo, K. Y., Shen, 2. Q., Zhao, J. H. & Ho, P. T. P. 1998, ApJ, 508, L61 Melia, F. & Falcke, H. 2001, ARA&A, 39,309 Pedlar, A,, Anantharamaiah, K. R., Ekers, R. D., et al. 1989, ApJ, 342, 769 Reid, M. J., Readhead, A. C. S., Vermeulen, R. C., & Treuhaft, R. N. 1999, ApJ, 524,816 Roberts, D. A. & Goss, W. M. 1993, ApJS, 86, 133 Yusef-Zadeh, F. 1989, in IAU Symp. 136: The Center of the Galaxy, 243 Yusef-Zadeh, F. & Wardle, M. 1993, ApJ, 405,584 Zhao, J., Bower, G. C., & Goss, W. M. 2001, ApJL, 547, L29
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Astron. NachrJAN 324, No. S1,397-401 (2003) I DO1 10.1002/asna.200385055
Closure Amplitude Analysis of 15,22 and 43 GHz VLBA Observations of Sagittarius A*: Size is Consistent with the Scattering Law G.C. Bower*’ I
Radio Astronomy Lab, UC Berkeley, Berkeley CA 94720
Abstract. High frequency very long baseline interferometry is necessary to avoid the effects of strong interstellar scattering when imaging Sagittarius A*. Unfortunately, the reliability of imaging declines with increasing frequency due to amplitude calibration errors for this low declination source. Closure amplitude analysis provides an unbiased method for estimating source parameters. We analyzed several new and several archival VLBA data sets with the closure amplitude technique. Our results indicate that there is no deviation from the scattering law in the minor axis size at 15, 22 and 43 GHz. There is a slight excess (20 3%) over the scattering law in the major axis size at 43 GHz which can be accounted for by a marginal increase in the error estimate, a slight recalibration of the scattering law or an intrinsic source of size 0.2 mas 2 AU 40R,. The absence of an apparent jet or outflow allows us to set an upper limit to the velocity of any ballistic components at 0.001~. N
N
N
1 The Size of Sagittarius A* Measurement of the size of the compact nonthermal radio source in the Galactic Center, Sagittarius A*, has a long and colorful history (Goss 2003). The discovery that the image of Sgr A* is an ellipse with major and minor axes that scale as X2, indicative of interstellar scattering, has pushed observations to short wavelengths where the scattering effects are weakest. The expectation is that intrinsic source structure will dominate the scattering effects at a short enough wavelength. This will allow us to separate the many different models for the emission of Sgr A* based on size and morphology. Lo et al. (1998) showed simultaneous VLBA observations that followed the scattering law from 6 cm to 7 mm wavelength in the major axis, aligned roughly East-West. On the other hand, the minor axis, aligned roughly North-South, showed an apparent increase above the scattering law at 7 mm. Combined with previous measurements by Bower & Backer (1998), Lo et al. inferred a > 40 detection of an intrinsic source with a size of 72R,. Unfortunately, there are numerous technical problems with making this measurement (Bower et al. 1999), principally related to the difficulty of accurately calibrating amplitudes at these wavelengths and at the low elevation of Sgr A* as viewed from North America. Doeleman et al. (2001) showed that the use of closure amplitude could be used to strongly and convincingly constrain the size of Sgr A* with VLBI observations at 3.4 mm. The closure amplitude is an amplitude gain independent quantity calculated for four stations m.,n,p and 4:
(1) where IV,, I is the amplitude of the visibility on the baseline between stations i and j . The closure amplitude is independent of all station-dependent amplitude errors, such as pointing errors, dish deformation, and variable opacity. It does not eliminate baseline-dependent errors such as variable decorrelation (which also influence conventional calibration and imaging techniques). The most significant loss from closure * Correspondingauthor: e-mail: gbowerOastro.berkeley.edu, Phone: 4 1 510 642 4075, Fax: +01 510642 4075
@ 2003 WILEY-VCH Verlag CmhH & Co KGaA. Weinhem
G. Bower: VLBA Observations of Sgr A*
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Radio Light Curve and New VLBA Observations 1.6-
T
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2001.5 2002 Time (years)
. .. . . . .
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Fig. 1 Light curve of Sgr A* at 43 GHz with VLBA observation dates indicated with red vertical bars. The first three observations were at 43 GHz and 22 GHz, while the later observations were at 43 GHz only.
amplitude analysis is the reduction in the number of degrees over freedom relative to a calibrated data set. The number of independent data points for a 7-station VLBA experiment is reduced by a factor 14/21. But this loss is more than offset by the confidence that the result gives through its accurate handling of calibration errors.
2 VLBA Observations and Analysis A number of new VLBA observations were obtained as a counterpart of the VLA radio flux density monitoring program of Sgr A* (Zhao et al. 2003, Figure 1). In addition to these observations a number of past VLBA experiments at frequencies of 15,22 and 43 GHz were reanalyzed, including those of Lo et al. Data were fringe-fit with AIPS first and then transferred to MATLAB and analyzed with proprietary code. The code forms the closure amplitude from the visibilities, averages the closure amplitudes and uses the scatter in the average to determine the error in the closure amplitude, and then uses a nonlinear fitting method for modeling the closure amplitude. Noise is added to the model closure amplitude, avoiding the problem of unbiasing the measured closure amplitude. A typical set of closure amplitude data points and models are shown in Figure 2. Errors in the model were determined by calculating x2 for a grid of models surrounding the solution and fitting constant x2 surfaces (Figure 3). Closure phases were also computed and analyzed. Non-axisymmetric structure is not reliably detected on any triangle, confirming the hypothesis that Sgr A* consists of only a single axisymmetric component. This is in spite of the range of flux densities and activity states observed.
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3 Closure Amplitude Results We show our results for 3 22-GHz and 7 43-GHz experiments in Figures 4 and 5. At 22 GHz, our mean result is consistent with the measurement of Lo et a1 and with the scattering size. This demonstrates the accuracy of both methods where calibration is more routine. The results at 43 GHz are more precise than at 22 GHz because of the greater fraction of the array contributing to the result: the source is smaller with respect to the array synthesized beam with increasing frequency. The 43 GHz minor axis size in individual measurements and in the mean result is significantly more compact than the extended source measured by Lo et al. The plot includes the data used by Lo et al. and reanalyzed with our procedures. This indicates that the North-South extension previously reported is not present and that intrinsic structure is
Mil
M. Pessah and F. Melia: Diffuse X-rays from Sgr A*
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d
I
0 m
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-2
-3 0.001
0.1 1 Distance from Sgr A‘ [pc]
0.01
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Fig. 3 Plot of the enclosed mass versus distance from Sgr A* for two extreme models considered in this paper. The dashed curve is a Plummer model for the visible matter. Open and filled circles correspond to total mass and inferred dark matter. The thick solid curve corresponds to the best fit model shown in Figure 1 ( B = 5 x lo1‘ G). The dotted curve corresponds to a similar model except for the lower bound adopted for the magnetic field ( B= 5 x l O I 4 G).
0.5
1 .o
1.5 Mbh
2.0
2.5
[ 1 06Ma1
Fig. 4 Contour plot for the enclosed mass versus radius on the r o - h f b h plane. Each curve is labeled by its specific x2 value, ranging from 0.1 to 3.0. In addition, the two points corresponding to the theoretical curves in Figure 3 are indicated by the circle (the best fit model) and the square (correspondingto the lowest value of h f b h for which a reasonable fit to the diffuse emission could be obtained).
1032-1033erg scl and the flux profile was lower on average. Note that the value of X2/dof is about 0.2 for the case with B 2i 1015 G, while it is around 1 for B N 1014 G . These results have several important implications for our understanding of the black-hole nature of Sgr A*. Clearly, this model does not provide an alternative to the black-hole scenario. On the contrary, it requires the presence of a massive point source to create a deep potential and provide the physical conditions for efficient accretion. However, it is no longer clear that Mbh = 2.6 x lo6 Ma.In fact, if the diffuse X-rays are produced by the cluster described here, our results point to an upper bound for the black-hole mass of 2.2 x lo6 M a . This value is in agreement with independent lower values obtained by numerical simulations of the dynamics of a massive black hole under the gravitational influence of a dense cluster. For example, Chatterjee et al. (2002) derived a lower bound of 1.1 x lo6 M o for the black hole at the Galactic Center. In this regard, it may be worth recalculating the emissivity of Sgr A* in models that depend rather sensitively on the black-hole mass. We have assumed that the cluster contains only neutron stars and have neglected any possible representation from white dwarfs and solar mass black holes. While these are unlikely to contribute significantly to the X-ray emission (see Haller et al. 1996), they may be important dynamically. We have also assumed that a fraction (1 0 %) of the neutron stars have strong magnetic fields that have not decayed below 1014 G over a Hubble time. Deviations from these assumptions should be reflected in the inferred mass of the dark cluster through the factor 6 in Eq. (2). It is also worth noting that our model implies that roughly 4 x lo3 strongly magnetized neutron stars are present within 1” of the supermassive black hole. If any of these erg s-’), the inferred luminosity were young neutron stars powered by magnetic field decay ( L x > would be inconsistent with current observations. However, star formation at the Galactic Center occurs in bursts, the latest of which peaked some 10-100 million years ago. It is therefore reasonable to assume that the last generation of magnetars faded long ago. Considering the simple model proposed here for the dynamical structure of the cluster, the overall agreement with the observed diffuse emission is reasonably good. The model appears to be viable and warrants additional, detailed study. Theoretically, the dynamics of such a cluster needs to be better understood, particularly with regard to its stability. In this work we set the parameter in the 7 model to 2.5. This value
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is physically motivated because it produces the internal slope of -1/2 in the density profile observed in earlier numerical simulations. Nonetheless, to test the sensitivity of our results on 77, we repeated our calculations with 71 = 2 (internal slope of -1). The results showed that it is possible to obtain an acceptable description of the actual flux distribution but the fit was not as good as the one shown in Figure 1. The mass of the cluster suggested by the best fit model for 77 = 2 was similar to the one found for 17 = 2.5, but the core radius was this time E 0.1 pc. Clearly, obtaining a more refined stellar density profile will reflect on the ultimate distribution of the diffuse emission. Also, the emergent spectrum from each star needs to be taken into account with greater care, since this will modify the number of stars required to match the observed luminosity. Finally, on the observational side, it would be desirable to acquire the spectral component of the diffuse emission separately from the point-source emission. Acknowledgements We would like to thank Daniel Eisenstein, Dennis Zaritsky and Scott Tremaine for useful discussions. We are also very grateful to Fred Baganoff and Mark Moms for providing the data prior to publication. This research was partially supported by NASA under grants NAG54239 and NAG5-9205, and has made use of NASA's Astrophysics Data System Abstract Service. FM is grateful to the University of Melbourne for its support (through a Miegunyah Fellowship).
References Binney J. & Tremaine S. 1987, Galactic Dynamics (Princeton Univ. Press, Princeton). Baganoff F. K., Maeda Y., Moms M. et al. 2002, ApJ, in press (astro-ph/0102151) Chatterjee P., Hernquist L. & Loeb A. 2002, ApJ, 572.37 1 Coker, R.F. & Melia, F. 1997, ApJ, 488, L149 Duncan, R. &Thompson, C. 1992, ApJ, 392, L9 Genzel R., Thatte N., Krabbe A,, Kroker H. & Tacconi-Garman L. 1996, ApJ, 472, 153 Genzel R., Eckart A,, Ott T. & Eisenhauer F. 1997, MNRAS, 291, 219 Ghez A,, Klein B., Morris M. & Becklin E. 1998, ApJ 509,678 Haller J., Rieke M., Rieke G., Tamblyn P., Close L. & Melia F. 1996, ApJ, 456, 194 Hey1 J.S. & Kulkarni S.R. 1998, ApJ, 506, L61 Kouveliotou, C. et al. 1994, in: AIP Conf. Proc. 307, Second Huntsville Gamma-Ray Burst Workshop, edited by G. Fishman, J. Brainerd & K. Hurley (New York ATP), 167 Kouveliotou, C. et al. 1998, Nature, 393,235 Maoz E., 1998, ApJ, 494, L181 Melia, F. 1992, ApJ, 387, L2.5 Melia, F. 1994, ApJ ,577, 426 Melia, F. & Falcke, H. 2001, ARA&A, 39, 309 Morris M. 1993, ApJ, 408,496 Murphy B. W., Cohn H. N. & Durisen R. H. 1991, ApJ, 360,60 Quataert E. 2002, ApJ, 575, 855 Quinlan G. & Shapiro S. 1990, ApJ, 356,483 Ruffert, M. & Melia, F. 1994, A&A, 288, L29 Rutledge, R. 2001, ApJ, 553,796 Tremaine S., Richstone D. O., Byun Y.I., Dressler A., Faber S. M., Grillmair C., Kormendi J. & Lauer T. R. 1994, AJ, 107,634 Zampieri, L., Turolla, R., Zane, S. & Treves A. 1995, ApJ 439, 849
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Astron. Nachr./AN 324. No. S I , 475 -48 I (20031 / DO1 10.1002/asna.200385045
A Relativistic Disk in Sagittarius A*
’
Siming Liu * and Fulvio Melia’ * 2
’ Center for Space Science and Astrophysics, Stanford University, Stanford, CA, 94305-4060
’ Department of Physics and Steward Observatory, The University of Arizona, Tucson, AZ,85721 Key words Supermassive Black Hole, Galactic Center, Radio Astronomy. PACS 04A2.5 The detection of a mm/Sub-mm “bump” in Sgr A*’s radio spectrum suggests that at least a portion of its overall emission is produced within a compact accretion disk. This inference is strengthened by observations of strong Linear polarization (at the 10 percent level) within this bump. No linear polarization has been detected yet at other wavelengths. Given that radiation from this source is produced on progressively smaller spatial scales with increasing frequency, the mm/Sub-mm bump apparently arises within a mere handful of Schwarzschild radii of the black hole. We have found that a small (10-Schwarzschild-radii) magnetized accretion disk can not only account for the spectral bump via thermal synchrotron processes, but that it can also reproduce the corresponding polarimetric results. In addition, the quiescent X-ray emission appears to be associated with synchrotron self-Comptonization, while X-ray flares detected from Sgr A* may be induced by a sudden enhancement of accretion through this disk. The hardening of the flare-state X-ray spectrum appears to favor thermal bremsstrahlung as the dominant X-ray emission mechanism during the transient event. This picture predicts correlations among the mm, IR, and X-ray flux densities, that appear to be consistent with recent multi-wavelength observations. Further evidence for such a disk in Sgr A* is provided by its radio variability. Recent monitoring of Sgr A* at cm and mm wavelengths suggests that a spectral break is manifested at 3 mm during cm/Sub-mm flares. The flat cm spectrum, combined with a weak X-ray flux in the quiescent state, rules out models in which the radio emission is produced by thermal synchrotron process in a bounded plasma. One possibility is that nonthermal particles may be produced when the large scale quasi-spherical inflow circularizes and settles down into the small accretion disk. Dissipation of kinetic energy associated with radial motion may lead to particle acceleration in shocks or via magnetic reconnection. On the other hand, the identification of a 106-day cycle in S g r A*’s radio variability may signal a precession of the disk around a spinning black hole. The disk‘s characteristics imply rigid-body rotation, so the long precession period is indicative of a small black-hole $pin with a spin parameter a / M around 0.1. It is interesting to note that such a small value of a / M would be favored if the nonthermal portion of Sgr A*’s spectrum is powered by a BlandfordZnajek type of process; in this situation, the observed luminosity would correspond to an outer disk radius of about 30 Schwarzschild radii. This disk structure is consistent with earlier hydrodynamical and recent MHD simulations and is implied by Sgr A*’s mm/Sub-mm spectral and polarimetric characteristics. For the disk to precess with such a long (106-day) period, the angular momentum flux flowing through it must be sufficiently small that any modulation of the total angular momentum is mostly due to its coupling with the black-hole spin. This requires that the torque exerted on the inner boundary of the disk via magnetic stresses is close to the angular momentum accretion rate associated with the infalling gas. Significant heating at the inner edge of the disk then leaves the gas marginally bounded near the black hole. A strong wind from the central region may ensue and produce a scaled down version of relativistic (possibly magnetized) jets in AGNs.
* Corresponding author: e-mail:
[email protected], Phone: +01650 723 01 12, Fax: 4 1 650 723 4840
* * Miegunyah Fellow
@ 2003 WILEY-VCH Verlag GmhH & Co K G A Weinheim
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1 Introduction The compact radio source, Sgr A*, at the dynamical center of our Milky Way Galaxy, is believed to be associated with a supermassive black hole (Melia & Falcke 2001). Evidence in support of this is quite compelling, especially with the detection of a 3 hour X-ray flare (Baganoff et al. 2001) from the direction of Sgr A* and the recent monitoring of a star orbiting within light-days of the black hole, which points to a central mass of 3.7 f 1.5 x lo6 Ma (Schodel et a1 2002), consistent with an earlier measurements of 2.6 x lo6 A4, (Eckart & Genzel 1996; Ghez et al. 1998). The nature of Sgr A* thus bears critically on our understanding of black-hole physics. Several mechanisms have been proposed over the past decade to explain its broad band spectrum and polarization, among them Bondi-Hoyle accretion (Melia 1992); a two-temperature, viscous disk (Narayan, Yi & Mahadevan 1995); a relativistic nozzle (Falcke & Markoff 2000) and a recent combination of advection-dominated disk with a nozzle (Yuan, Markoff & Falcke 2002). Over the past decade or so, our group has adopted a theoretically-motivated phenomenological approach (Melia, Liu & Coker 2000,2001 j, in which the observations play a crucial role in constraining the theoretical picture. The detection and confirmation of a mm/Sub-mm bump in Sgr A*'s spectrum suggests an emission component different from that responsible for the cm radio emission (Zylka, Mezger & Lesch 1992; Falcke et al. 1998). This emission component is also implied by Sgr A*'s variability. Radio Observations show that Sgr A*'s fluctuation amplitude increases toward high frequency (Zhao & Goss 1993) and there is a spectral break at 3 mm during radio flares (Zhao et al. 2003). Since high-frequency radio emission is produced by relatively more energetic particles, located deeper in the gravitational well of the black hole (Melia, Jokipii & Narayanan 1992), the mm/Sub-mm emitting gas should be very close to the black hole's event horizon. The detection of linear polarization in this spectral bump enhances this inference further and sets severe constraints on possible explanations for this component (Aitken et al. 2000). No significant linear polarization has yet been detected at frequencies lower than 112 GHz, though relatively strong circular polarization persists in the cm band (Bower et al. 2002). The flip of the position angle of the polarization vector by about 90" between 230 GHz and 350 GHz favors a scenario where the mm/Sub-mm emission is produced within a small, optically thin, magnetized accretion disk (Melia et al. 2000). No other model so far can explain this linear polarization characteristic (In the empirical model of Ago1 (2000), the frequency where the position angle flips by 90" is much lower than the frequency corresponding to the spectral peak of the flux density, which is not in line with the observations). The existence of a small disk is also motivated theoretically. Earlier hydrodynamical simulations suggested that black-hole accretion from stellar winds, as is the case for Sgr A*, is characterized by a small angular momentum of the captured gas (Coker & Melia 1997). This accreted angular momentum is too small for the gas to settle onto a large disk, as required by the ADAF model (Narayan et al. 1995). The captured angular momentum is instead barely sufficient to circularize the gas just before it falls across the black hole's event horizon. Detailed MHD simulations have provided an indication of the structure for such a disk (Hawley & Balbus 2002). The fact that the Magneto-Rotational Instability (Balbus & Hawley 1991) can induce a MHD dynamo in the disk provides a straightforward explanation for the mm/Sub-mm bump as the result of synchrotron process. The magnetic field also provides an anomalous viscosity and, given its strength, couples the electrons and ions via reconnection, so that a single temperature fluid is maintained (Melia et al. 2001j. X-ray observations of Sgr A* have provided additional means of learning about its nature. It turns out erg s-l that Sgr A* is an extremely weak X-ray source with a quiescent X-ray luminosity of 2.2rE x (Baganoff et al. 2001j. Interestingly, electrons responsible for the mm/Sub-mm emission also Comptonize the radio photons into the X-ray band. It is notable that the physical conditions required to produce the mm/Sub-mm spectrum can also account for Sgr A*'s quiescent X-ray emission. Chundvu also detected a strong X-ray flare from the direction of Sgr A*. The flare lasted about 3 hours and featured a variation on a 10 minute time scale, suggesting an emission region no bigger than 20 T S , N
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Astron. Nachr./AN 324.No. SI (2003)
where T S = 7.7 x 1011 cm is the Schwarzschild radius for the 2.6 x lo6 111,supermassive black hole associated with Sgr A*. The peak flux density for this flare is 50 times higher than that in the quiescent state. Recent X-ray observations have shown that this type of X-ray flare is common to Sgr A*, occurring about once per day (Baganoff et al. 2003; Goldwurm et al. 2003). The fact that the flares are very strong, are variable on a 10 minute time scale and have a flat spectrum, pose telling theoretical challenges which, at the same time, also create a valuable opportunity for constraining the physical conditions near the black hole’s event horizon. Our study has shown that an enhancement of the mass accretion rate through the disk can not only account for these flares, but can also induce strong Sub-mm/Far-IR flares that should occur simultaneously with the X-ray flares (Liu & Melia 2002a). Recent observations of Sgr A*’s mm/Sub-mm variability has indicated that the Sub-mm spectral index increases significantly during radio flares (Zhao et al. 2003), consistent with our prediction. However, radio flares usually last for several days, which is much longer than the duration of an X-ray flare, suggesting more complicated physical processes. Nevertheless, the nondetection of an IR flare (Hornstein et al. 2002) seems to favor this model, where X-ray flares are produced via thermal bremsstrahlung processes, over the nozzle model, in which synchrotron self-Comptonization is introduced to account for the X-ray flare emission (Markoff et al. 2001). Moreover, the low quiescent X-ray flux also delimits hot gas content around Sgr A*. When this constraint is combined with the flat radio spectrum, one can show that the cm radio emission from Sgr A* cannot be produced by a bounded, thermal synchrotron source (Liu & Melia 2001). One possibility is that the radio emission is produced via nonthermal synchrotron processes in the region where the large scale quasi-spherical inflow circularizes to form the small accretion disk responsible for the mm/Sub-mm and X-ray emission. Energetic, nonthermal electrons can in principle be produced by the dissipation of kinetic energy associated with the radial motion of the infalling gas in shocks or magnetic reconnection. Assuming that a fixed fraction of particles is accelerated in this way, one can obtain a good fit to the radio spectrum. The circular polarization properties may then be associated with the turbulent nature of the gas in this region (see, e.g., Beckert & Falcke 2002; Ruszkowski & Begelman 2002). On the other hand, a 106 day period in Sgr A* radio variability recently reported by Zhao et al. (2001) appears to be associated with the precession of a small hot disk under the influence of a spinning black hole (Liu & Melia 2002b). The physical characteristics of the disk indicate that it will precess as a rigid-body. However, for the disk to survive longer than the observed period, the net angular momentum flux through the disk must be extremely small, which requires that the inward angular momentum flux associated with the accreting gas must be cancelled almost completely by the outward angular momentum induced by torque associated with the magnetic stresses. A nonzero torque at the inner edge of an accretion disk has been discussed extensively (Krolik 1999; Gammie 1999; Ago1 & Krolik 2000) during the past few years. Recent MHD simulations have also confirmed several of these theoretical speculations (Hawley & Balbus 2002). Should this picture be correct, it should be noted that a small black hole spin of 0.1 M , where M is the mass of the black hole, would be favored if the nonthermal portion of Sgr A*’s spectrum is instead powered by energy extracted from the black hole via a Blandford-Znajek type of process. The 3Ors, consistent with the general precession period then requires that the disk has an outer radius of picture described above. The power extracted from the black hole also heats up the gas near the event horizon and unbinds it. The ensuing wind is not unlike the relativistic jets observed in AGNs. Further exploration of this idea may eventually reveal a more refined view of the processes hidden in the central engine of these sources.
-
N
2 A Relativistic Disk Model for the mm/Sub-mm Emission from Sgr A* The model of a hot, magnetized, small accretion disk in Sgr A* has been developed fully in the paper by Melia et al. (2001) where, prior to the availability of all the observational constraints described above, the inner boundary condition was chosen to have zero torque. In this instance, the temperature at the outer boundary of the Keplerian region is the primary free parameter. The disk structure is determined once
47 8
S. Liu and F. Melia: Relativistic Disk
one specifies the inner (rJ and outer ( r o )radii, the magnetic (&) and viscous (By)parameters, the mass accretion rate hk and the inclination angle of the disk. The best fit to the mm/Sub-mm polarization and spectral data is shown in Figure 1 (Melia et al. 2000). Note that here a negative percentage means that the position angle of the polarization vector is parallel to the angular momentum vector of the disk, while positive polarization means that the polarization vector flips by 90" with respect to negative polarization. 3.3 x 10l1 Hz, which is higher than the peak The frequency at which the polarization vector flips is frequency of the flux density, 2.1 x lo1' Hz. This is a unique feature of our relativistic disk model, that is apparently not yet matched by alternative scenarios (cf. Ago1 2002). It is straightforward to understand these polarization characteristics. At mm wavelengths, the red shift side of the disk becomes optically thin first. At this point, the emission is mostly from the front and back of the black hole, where it is polarized in the direction parallel to the disk's spin axis due to the influence of the very strong toroidal field within the disk. At Sub-mm frequencies, even the gas to the front and back of the black hole becomes optically thin, and the emission from the blue shifted side of the disk dominates; the polarization vector thus flips by 90". Faraday rotation by the intervening plasma will make the observed flip of the polarization vector different from go", which can reconcile the slight difference between the theoretical prediction and the observational results. Due to the relatively poor angular resolution of JCMT (22" at 220 GHz), the corresponding error bars are. quite big, as can be seen from Figure 1 (Aitken et al. 2000). However, the detection of strong linear polarization is quite obvious. Recent high resolution (3.6" x 0.9") BIMA observations have confirmed strong linear polarization at 220 GHz (Bower et al. 2003), adding some confidence to the model. N
N
4
1
I
,
,
,
I
.
,
,
.
.
,
.
RIUA JCW
0
-.,
-2
\
r '
I
-4
5 -6
-8
-10
Fig. 1 Best fit to the linear polarization of radio emission from Sgr A*. Here the model parameters are as follows: M = 4.1 x 1OI6g s-', PP = 0.02, Py = 0.2, ri = 1.8rs and r , = 8.5 r s . The inclination angle of the disk is 30". At the outer boundary, the gas temperature is fixed by the assumption that the thermal energy of the gas equals 7% of its dissipated gravitational energy.
Fig. 2 Best fit to Sgr A*'s quiescent spectrum. The model parameters are shown in the figure. Here the disk has an inclination angle of 45". The thermal energy of the gas is assumed to equal 80% of its dissipated gravitational energy at ro. The dashed line here denotes emission from the small disk. The dotted line gives emission produced by nonthermal particles in the circularization zone.
3 X-ray Emission from the Relativistic Disk Chandra observations indicate that quiescent X-ray emission from Sgr A* is very weak and soft, with an X-ray luminosity 2.2-t::; x erg s-l and a spectral index 1.5:$, which is not consistent with an ADAF (Baganoff et al. 2001). Given the fact that the plasma is so hot in the disk that electrons are relativistic
Astron. Nachr./AN 324, No. S 1 (2003)
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(Melia et al. 2001), one is motivated to investigate the effects of synchrotron self-Comptonization (SSC) in this medium. Radio variability observations and recent theoretical developments (e.g., Liu & Melia 2002b) favor a nonzero stress at the inner edge of the accretion disk. We will henceforth adopt a zero angular = -P,&R,T/prR. The momentum flux condition. The radial velocity of the accretion flow is then IJ, other equations derived in Melia et al. (2001) are still applicable and we won't reproduce them here. In Figure 2, we provide the best fit to Sgr A*'s broadband spectrum. SSC evidently accounts for Sgr A*'s quiescent X-ray spectrum very well.
2
,x
0
-2
\
-.'
-zs
-4
-6
-a 2
--. -
0
-2
2
-4
-b
-6
0
-8 -10
I
...
I . . . I . . . I . . . l . . . -
Fig. 3 An accretion induced X-ray flare from Sgr A*. The righthand panels give the temperature (thin lines) and density profiles (thick lines) for the disk ( T ~= r s / 2 ) .The left panels show the corresponding disk spectra. Note that the model prediction is consistent with the IR upper linuts (Hornstein et al. 2002).
However, Sgr A*'s X-ray emission during a flare is much more complicated than that during quiescence. Although the short variation time scale of 10 minutes is consistent with the flare being induced by an accretion process in the disk, the fact that the X-ray flux density can increase by a factor of 50 suggests dramatic changes in the disk's structure. Moreover, the hardening of the X-ray spectral index also rules out SSC as the dominant X-ray emitting mechanism during the flare (Liu & Melia 2002a). We note that the hot disk described here is not stable when the mass accretion rate is large. When A? increases, bremsstrahlung cooling becomes more and more important and can be the dominant cooling mechanism. The X-ray flare may in fact be associated with enhanced thermal bremsstrahlung emission. Figure 3 depicts such a scenario. Here the disk has an outer radius of 9 rs and we assume that the enhancement of accretion through it can suppress the MHD dynamo. A justification for this is that cooling becomes more efficient with increasing Ak, and this decreases the gas temperature. We infer that & is thus anti-correlated with A?. During the X-ray flare, the gas density can be as high as 10'' cmP3 and the magnetic field reaches 100 Gauss. The corresponding synchrotron cooling time is a few hours, consistent with the general picture outlined above (Petrosian 1985). Recent multi-wavelength observations have indicated that there is no obvious flux change at 3 mm during the X-ray flare (Baganoff et al. 2003). According to our model, the disk is optically thick at 3 mm, so the flux density is not expected to change significantly (see Figure 3).
4 The Nature of Radio Emission from Sgr A* In Figure 2 we shiiwcd a fit tii thc cm cmissiun from Sgr A* under the a s s u r n p h lhat this radiation is produccd via nonthcrinal synchrotrrin prrmsscs in thc circularir,ation zonc (the Inodcl ddails can he f'uund in I,iu & Mclia 2001 ). The dintritutirin ( i f nonthennal Iurticles i s givcn hy N(E.r.) : 1 . 7 x 10 I:' - " ' 7 ~ ( r )whcrc % E is the electron encrgy, and 71,(r)i s thc cleclron density at radius T . Althuugh magnetic rcctinnection at smaller radii of the disk can also induuc pitrticlc accelcraliun (adding to the contribufiun made by (he nonihcmal parliclcs in thc circul;ir,zation zone), the fact h a t the gas tempr.raturc is as hiph as 10 M c v therc appcars to make thc thcrrnal proccss dominarit. .UevcrhlcPs, ii complete
treatmcnt of this pruhlctn i n c o q i o r k n g particle acceleraticlrr via magnetic 1,econncctiori is warranled. Vie can understarid the nnnthcrinal natiIrc uf cin radio emission using Ihc following argument. Hccausc quiesccnt X-+ly emission from Sgr A" i s w a k (sec Fiigurc 2)- H'C can constrain thc hut gas contcrit in Sgr A + via its hremsslrahlung emissivity. If u'c a.wrIic That thc gas is bounded, h e gas lcmpcraturc musL he lower than its virial value. Cornhining thcsc two upper limits, one can show that to produce the 1.36:[.;Hi: flux from Sgr A * rhc miignctic field energy density musl he moxc than ten times higgcr than thc thermal cncrgy densit) of the hoi pas. Such ii configuration is no1 physical iT the rrisgnctic field is intrinsic to thc hot gas. Of coursc, it is iilso prissihlc lhal the radio einisqion is produccd by m n c unhonitdcd plaxma, as prtiposcd iri the Jet model hy Markoff arid Palckc (2000). Then t h e origin of the jet becomes !he largc unknown i n thc rririrlcl. Tlic dctecliun o f a 106 d;iy radio cycle is intriguing hecauye i L is intrinsic lu Sgr A" (%hat> ti al. 20(11) and rcccnt V I A o1)aervatiuris in& catc [hat ihe erriission is prnduccd x i t h i n 1,40r,! (Hoacr ct al 20132). Thc dynainical t i m e scalt. within such a sniall region is muck sburlcr than this period, suggesting il inay he as.m%led with ail iniriii 0). The vertical scale height H is introduced through the hydrostatic balance as H = cs/O where 0 is the angular velocity. Since the latter is assumed to be Keplerian, the radial momentum equation (eq. 2 in NY94) is trivially satisfied. Since the cold disk is also Keplerian, there is no exchange of specific angular momentum between the two flows and the angular momentum conservation equation (3 in NY94) is unaltered. With at = O K we get
it?
127ra d [pc2R2H] , RRKd R
= 4rRHp.u~ = __-
where Q is the viscosity parameter (Shakura & Sunyaev 1973). The entropy equation should include (in addition to the usual terms) the thermal conduction flux, F,,, and the hydrodynamical flux of energy in the vertical direction. To derive this equation, we follow the formalism of Meyer & Meyer-Hofmeister ( 1 994; see their equation 8), with the following exceptions. We neglect winds here because we are interested in cooler, condensing solutions. Thus their sideways term (their eq. 5 ) is not included. In addition, the radial entropy flow term (the “advective cooling”; NY94) is designated Qadv. Following NY94 we set Qadv = fadv&+, where fad,, 5 1 is a parameter. Further, for subsonic flows 0, then the hot flow is condensing onto the cold disk. In the opposite case of a small (p and a “large” a, b < 0 , viscous heating prevails. Thermal conduction then serves to evaporate the inactive disk. We should also note that the regime of large a ( w 0.3) was already studied by F. Meyer and collaborators in many papers. Their solutions are for higher accretion rates and therefore they are in the non-saturated regime, which is roughly speaking equivalent to the 4 > a. Inserting now 11, = -bc, into equation (1) and also substituting U R for its value found from equation (2), we arrive at a second order differential equation that contains two variables, p and c,:
bpRc, = 3 a -
a dR
1 8 {[pc:$]} ROK dR -
.
This equation cannot be solved in a general case. By introducing ZJ, # 0 we added an extra variable to the accretion flow equations, and the number of independent equations is now smaller than the number of
S. Nayakshin: Inactive Disk in Sgr A*
486
unknowns. (In particular, both v, and c, are to be found from the single energy equation 3). This situation is well known in analytical ADIOS wind solutions (BB99). In the latter case one has to introduce three free parameters that describe the mass, energy and angular momentum carried away by the wind. The most natural way to proceed here is to suggest that the temperature is a power-law function of radius. For example, for AJlm, T ( R )0: R-‘ in a broad range of radii. On the other hand, thermal conduction tends to smooth out temperature gradients (for example within supernova remnants), and hence in the other extreme T ( R )N const (we in fact observed this nearly constant T ( R )in our numerical simulations). Thus, cs = co(R,/R)’, where 0 5 X 5 1/2 and co is the sound speed at R,. If we define u = then equation (5) can be re-written as
a,
Finally, defining i j a2fi
-=-au2
= pu7-6x,we obtain
1 12
fi y4(1-’)
(7)
’
Here 1, is the dimensionless “condensation length”:
If there were no thermal conduction losses, the internal energy of the hot gas would be of order its gravitational energy. If the inactive disk extends to R > R,, the thermal conduction will reduce the gas thermal energy. Hence we expect that czR,/GMBH 5 1. We are most interested in the case 4 >> a and therefore we shall only explore the X = 0 case below.’ a/q5 1. Let us now explicitly write down the results from this approximate solution:
Here p, is gas density at R,, and A& is the accretion rate at that point. Note that equation (1 1) shows that the radial velocity is constant and is substantially smaller than the sound speed. Further, for 1, < 1/2 (recall that we assumed 1, > a w e h a v e C = l and 1, > 3R, because we neglected any relativistic corrections to the gravitational potential. In addition, the derived expression is only a rough R (which estimate since we assumed vertically averaged equations which are clearly inaccurate for H in our simple model occurs at R = R,). One can in fact show that realistically &hol should be smaller by
488
S . Nayakshin: Inactive Disk in Sgr A*
a factor of at least few. The point here is that our simplistic solution is not bound at R = Rc because its thermal energy exceeds gravitational energy at that point. So either a wind takes away the excess energy (as BB99 argued for ADAF solutions) or more likely this simply means that we over-estimated the gas temperature at Re. We should also explicitly insert the disk inclination angle into the definition of the apparent efficiency since the observed luminosity of the thin disk is approximately proportional to cos i. Variation of the predicted spectrum with R, is shown in Figure (3). Summarizing, the apparent bolometric efficiency of the freezing flow is
4 Discussion We have suggested here that there exists an inactive disk around Sgr A*, a remnant of past powerful accretion (and probably star forming) activity. The disk may be quite light compared with both the BH and M a of the hot gas present in the the star cluster (Nayakshin et al. 2003), yet it easily out-weighs the region interior to the Bondi radius. The disk then serves as a very efficient cooling surface for the hot flow. The flow essentially gets frozen (stopped), and its energy is radiated as thermal emission at frequencies much below the X-ray band. The X-ray emitting flow thus simply disappears from “the radar screen”. This in our opinion may be the explanation of the exceptional apparent X-ray radiative inefficiency of Sgr A*. Further, since the hot flow does not penetrate very deep into the BH potential well, its total bolometric few R,/R, times smaller than that expected if the gas made it all the way into the BH luminosity is and was radiatively efficient. Thus the flow also appears radiatively inefficient in the bolometric sense. As we explained in the paper, we believe that the accretion of the winds from the hot stars is simply delayed in time and it is by no means radiatively inefficient in the long run. Falcke & Melia (1997; FM97) have assumed the hot wind infall as given and studied viscous evolution of the “fossil” disk on long time scales, whereas we concentrate on much shorter time scales on which the 106ff-l years for Tdl& = 100 K and R = lO4RY). structure of the disk does not change (i.e t < t,,,, Our study is therefore complimentary to that of FM97. N
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Our results concerning the conditions under which the wind-disk (or the hot flow-disk) interactions will not violate the tight NIR limits are quite similar to that of FM97. In particular, FM97 note that “the BondiHoyle wind must be accreting with a very high specific angular momentum to prevent it from circularizing in the inner disk region where its impact would be most noticeable”. We find that the circularization radius should be 2 3 x 104Rg,implying a very large angular momentum indeed. Further, we suggested that the disk and the hot flow angular momenta are at least approximately aligned or else there would be a substantial heating due to friction between the two, a heating not included in our analysis. It remains to be seen whether results will be qualitatively similar if the disk and the hot flow rotation axes are misaligned. FM97 considered a ‘‘large’’ value of R, being rather unlikely. We however note that according to recent data, the stars from which the hot wind originates appear to be on tangential orbits counter-rotating the few arcsecond few x 105Rgoff Sgr A* (e.g. Genzel 2000).There is thus Galactic rotation and are no deficit of angular momentum at these distances. Finally, we only studied here the region of the flow interior to R,. However the exchange of the angular momentum between the hot flow and the disk should take place at R > R, where the hot gas is sub-Keplerian. This should enrich the hot flow with the angular momentum directed as that of the the disk and hence the circularization of the hot flow should occur even if it had zero angular momentum at infinity. Thus the requirement of a large angular momentum in the wind may be relaxed, although this effect remains to be quantified with future calculations.
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In this paper we suggested that there exists a very cold inactive disk in Sgr A*, and that its role in the accretion picture is significant. While the hot gas is very tenuous and cannot radiate its energy away, it can easily transfer its energy into the cold disk via thermal conduction. The cold disk is much denser and much more massive than the hot flow and can serve as a very powerful freezer (or radiator) for the hot flow. As the hot flow looses its energy, it also looses its viscosity and “sticks” to the cold disk. The accretion flow is thus quenched by this seemingly “rzon-radiative”cooling. One can easily check that neither internal viscous dissipation nor the heat input from the hot flow in Sgr A* are sufficient to overcome the radiative cooling and restart the accretion in the inactive disk. It appears that only arrival of a new large supply of low angular momentum material could revive the inactive disk in Sgr A* now. We thank H. Falcke, C. McKee, F. Meyer, R. Narayan, R. Sunyaev and H. Spruit for discussions.
References Baganoff F. K., et al., 2001, Nature, 413, 45 Baganoff F. K., et al. 2003, these proceedings Blandford, R., & Begelman, M.C., 1999, MNRAS, 303, LI Cowie, L.L., & McKee, C.F. 1977, ApJ, 21 I , 135 Falcke, H, & Melia, F. 1997, ApJ, 479,740 Genzel, R. 2000, in the Proceedings of the Star2000 Meeting, editor R. Spurzem (astro-pW0008119) Hornstein, S.D. et al. 2003, these proceedings. Melia, F., & Falcke, H. 2001, ARA&A, 39, 309 Menou, K., & Quatdert, E. 2001, ApJ, 552,204 Meyer, F., & Meyer-Hofmeister, E. 1994, A&A, 288, 175 Miyoshi, M., et a]. 1995, Nature, 373, 127 Nardyan, R., & Yi, I. 1994, ApJ, 428, L13 Narayan, R., Yi, I., & Mahadevan, R. 1995, Nature, 374,623 Narayan R., 2002, p. 405, in “Lighthouses of the Universe”, Springer 2002, editors Gilfanov, M., Sunyaev, R., & Churazov E. Nayakshin, S., & Sunyaev, R. 2003, submitted to MNRAS (astro-ph/0302084) Nayakshin, S., et al. 2003, in preparation. Quataert, E., et al. 1999, ApJL, 525, L89 Quataert, E., 2003, these proceedings Schodel R., et al. 2002, Nature, 419, 694 Shakura, NJ., & Sunyaev, R.A. 1973, A&A, 24,337 Siemiginowska, A., Czerny, B., & Kostyunin, V. 1996, ApJ, 458,491
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Astron. Nachr./AN 324, No. S1,491-495 (2003) / DO1 10.1002/asna.200385046
Gamma-ray emission from an ADAF around a Kerr black hole Kazutaka Oka*' and Tadahiro Manmoto**2
' Department of Earth and Planetary Sciences, Kobe University, Kobe 657-8501, Japan
* Department of Physics, Chiba University, Chiba 263-8522, Japan
Key words accretion, accretion discs - black hole physics - Galaxy: center - gamma rays: theory PACS 04A25 We investigate the gamma-ray spectrum emitted from an ADAF and its dependence on the spin parameter of a central Kerf black hole, in order to examine whether the spectrum can be used to probe the spin parameter of black holes. We consider that the gamma-rays are produced through the decay of neutral pions created by proton-proton collisions in the vicinity of the central black hole. Since the energy distribution of the ion particles in an ADAF is not known, we consider two types of proton energy distributions: a thermal distribution and a power-law distribution. In the thermal model, we find that changes in the spin parameter from-0.95 to 0.95 can enhance the gamma-ray intensity by orders of magnitude. Thus, if the proton gas in an ADAF has a thermal distribution, the gamma-ray spectrum can be used as a probe to investigate the spin parameter of the central black hole. In the nonthermal model, on the other hand, the gamma-ray intensity is much less sensitive to the changes in the spin parameter than in the thermal model, and it would he difficult to estimate the spin parameter from the gamma-ray spectrum. We apply our model to the Galactic Center, Sgr A'. The unidentified gamma-ray source 3EG J17462851 is observed towards Sgr A* by EGRET. Our results show that the gamma-ray intensities predicted from our models are much lower than observations and we cannot find the spin parameter. We, however, consider that this is not a serious problem against our model since it is unclear whether the observed gammarays are from a point or a diffuse source at the Galactic Center. In order to investigate the spin parameter via the gamma-rays from the Galactic Center instruments with higher angular resolution is needed such as GLAST.
1 Introduction The dimensionless spin parameter a is an important physical quantity representing the black hole spin. If we could determine the spin parameter from observations, it provides not only the confirmation of the accretion theories, but also some insight into the accretion history of supermassive black holes in active galactic nuclei (AGNs). Then, how can we investigate the spin parameter of a black hole? In the present study we consider a Kerr black hole surrounded by an advection-dominatedaccretion flow (ADAF), which is a geometrically thick, optically thin hot accretion flow and has low radiative efficiency (Ichimaru 1977; Narayan & Yi 1994,1995a,b; Abramowicz et al. 1995). Emissions in radio to X-ray bands are determined by the cooling processes of electrons such as synchrotron, bremsstrahlung, and Compton processes. In the gamma-ray band, on the other hand, Mahadevan, Narayan, & Krolik (1997) pointed out that the ion temperature in an ADAF close to a black hole is high enough (Tt lo1' K) to produce gammarays through the decay of neutral pions, which are created by proton-proton (p-p) collisions. Their method is very interesting, because once the spectrum produced by the electrons is fixed, which means that all the N
' Corresponding author: e-mad: kazutakaakobe-u.ac.jp, Phone: +81 78 803 5754, Fax: +8178 803 5747 * * Corresponding author: e-mail: manmotoQ astroschiba-u.ac.jp @ 2003 WILEY-VCH Verlag GmbH & Ca KGaA, Weinhem
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parameters in an ADAF are determined, no additional parameters are required to calculate the gamma-ray spectrum. Thus one can calculate the gamma-ray spectrum uniquely. They, however, considered only the case of Schwarzschild black hole, and thus did not investigate the case of rotating black hole. Manmoto (ZOOO), on the other hand, studied the ADAF spectrum around a Kerr black hole by taking into account only the cooling processes of electrons. Manmoto (2000) showed that dependences of ion temperature and electron temperature on the spin parameter are different. Thus, if we combine the gamma-ray spectrum, which contains information on the ion temperature, and the spectrum in other bands, which contains information on the electron temperature, the spin parameter of a black hole could be determined. In the present study we examine whether the gamma-ray spectrum can be a probe to study the spin parameter. So far the ADAF model successfully explained spectra of several objects including the Galactic Center, Sgr A*. Among them, the Galactic Center is the only object whose gamma-ray intensity is above the detection threshold of EGRET (Mahadevan et al. 1997). In the present study we apply our model to the Galactic Center.
2 Description of Model 2.1 Gamma-ray Emission Mechanism In the ADAF model, the dissipated energy is stored in the accretion flow and advected inward. Since ions hardly radiate, they are heated almost up to the virial temperature. The ion temperature in the vicinity of a central black hole is so high (T, 10°K) that a p-p collision produces a neutral pion, T O , which then decays into two gamma-ray photons (Mahadevan et al. 1997). The emergent shape and luminosity of the gamma-ray spectrum from an ADAF dramatically depend on the proton energy distribution, which is in turn determined by the mechanism of viscous heating. However, the detailed mechanism of the heating is not well understood at present, and therefore it is not known whether the viscous heating leads to a thermal or a nonthermal distribution of proton energies (see Mahadevan & Quataert 1997). In order to calculate the gamma-ray spectrum we assume two different proton energy distributions: a thermal distribution (eg., Dermer 1986; Mahadevan et al. 1997) and a power-law distribution (Mahadevan et al. 1997; Mahadevan 1999). The details of the calculations of the gamma-ray spectrum are described in Oka & Manmoto (2003). N
2.2 Calculation of ADAF The structure of an ADAF is determined by the following parameters: the viscous parameter a, the ratio of the gas pressure to the total pressure &, the mass of the central black hole M , the mass accretion rate M , the fraction of the viscous heating that goes into electrons 6, and the spin parameter a(-1 < a < l), where a positive (negative) a means that the black hole corotates (counterrotates) with the accretion flow. In the calculation of the ADAF structure, we assume the ion particles to be thermalized, and therefore we obtain the ion temperature at each radius. For the nonthermal model we redistribute the energy of the ion particles with the total energy at each radius fixed. We assume that the electrons are always thermalized by action of self-absorbed synchrotron photons (Mahadevan & Quataert 1997). The details of the calculations of the ADAF structures are described in Manmoto (2000).
3 Results 3.1 Typical AGN-mass black hole We show the results obtained for the case of a typical AGN-mass (A4 = lO*M,) black hole. Three spin parameters of a = 0.95,0, and -0.95 are investigated. We assume that the spectra due to electron cooling processes have the same intensity at an X-ray point (1 keV). To do this, we adjust the mass accretion rate = 1 . 0 310-3A&, ~ h;r = 10-3&'c, and 5 . 8 10-*&fC, ~ for the models with a = -0.95,0, and 0.95 to be
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Fig. 1 Lef: Spectra in the thermal model for a = 0.95 (solid lines), 0 (dashed lines), and -0.95 (dotted lines). ~ For a = -0.95 (dotted Gamma-ray spectrum due to the pion-decay is located at frequencies higher than 0 - 1 0 ~Hz. lines), the actual gamma-ray part of the spectrum is 10 times lower than shown. Righr: Spectra in the power-law model with s = 2.75. Gamma-ray spectrum for a=0.95 (solid line) and a=O (dashed line) are very similar, and thus overlapped. In both figures the black dots show the X-ray point (1 keV).
respectively. Here Ak, is the Eddington mass accretion rate. Other parameters are set to be (Y = 0.1 and /?, = 0.5. We also assume that almost all the dissipated energy heats the ions by setting b = 0.001, although the determination of the value of 6 is still a controversial issue (see e.g., Bisnovatyi-Kogan & Lovelace 1997). Using the ADAF+Kerr model (Manmoto 2000), we obtain the structure of the flow, and then calculate the gamma-ray spectrum. Here the Doppler and gravitational shifts, or the bending of the photons path are not taken into account. The left panel of Figure 1 shows the spectra from the radio hand to the gamma-ray band for a = -0.95, 0, and 0.95 in the thermal model. The spectrum due to the pion-decay is located at frequencies loz1 Hz, while the spectrum due to the electron cooling processes such as synchrotron, higher than v bremsstrahlung, and Compton processes appears at frequencies lower than v N lo2' Hz. We can see that the gamma-ray intensity is enhanced by orders of magnitude when the spin parameter changes from -0.95 to 0.95. Based on the calculation, we can conclude that if proton gas in an ADAF has a thermal distribution, the gamma-ray spectrum can give a constraint on the spin parameter of the central black hole. The right panel of Figure 1 shows the spectra in the nonthermal model. We set the power-law index s of the power-law distribution to be 2.75. Other parameters are the same as in the case of the thermal model. We find that the gamma-ray intensity for each spin parameter has almost the same magnitude. Thus, it would be difficult to estimate the spin parameter from the gamma-ray spectrum. N
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Application to Sgr A*
We apply our model to the Galactic Center, Sgr A*. The spectrum of Sgr A* has been explained naturally with the ADAF model (Narayan, Yi & Mahadevan 1995; Manmoto, Mineshige, & Kusunose 1997; Narayan et al. 1998; Manmoto 2000; Narayan 2002). Although there are several objects whose spectra are explained by the ADAF model, Sgr A' is the only object whose gamma-ray intensity is above the detection threshold of EGRET (Mahadevan et al. 1997). Our models are adjusted so that the predicted emission agrees with the X-ray flux in quiescence measured by Chandra (Baganoff et al. 2001). We adopt both the thermal and the power-law distributions of proton energies. The model prediction is compared with the gamma-ray source 3EG 51746-2851 observed towards Sgr A* by EGRET. The ADAFparameters are as follows. The mass of Sgr A* is M = 2.5x106Ma,
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Fig. 2 Radio to gamma-ray spectrum for Sgr A*. Both thermal and power-law distributions of proton energies are adopted. In the gamma-ray band, the spin parameters are a = 0, -0.95, and 0.95 (from top to bottom) for the powerlaw model, and a = 0.95, 0, and -0.95 for the thermal model. For the thermal model with a = -0.95, the actual gamma-ray intensity is 100 times lower than shown. The gamma-ray data corresponds to 3EG 517462851 observed by EGRET (Hartman et al. 1999).
a is 0.1, pp is 0.5, and 6 is 0.1. Power-law index is 2.75. Spin parameters and mass accretion rates are (a, A?>= ( 0 . 9 5 , 1 . 7 4 ~ 1 0 ~ ~ A( k0 ~, 8) ., 6 ~ 1 0 - ~ A ? and ~ ) , (-0.95,l x10W5hj,). Figure 2 shows the radio to gamma-ray spectrum for Sgr A*. All the predicted gamma-ray intensities are lower than the EGRET data by more than two orders of magnitude, and therefore we cannot decide the spin parameter for Sgr A*. However, we do not consider that it is a serious problem against our model; our model suggests that most of the gamma-ray flux obtained by EGRET does not originate in Sgr A*. Since the angular resolution lo,it would have contamination from other sources such as radio arc and Sgr A" East of EGRET is (see e.g., Pohl 1997; Melia et al. 1998a,b; Markoff, Melia, & Sarcevic 1999). In order to remove such contamination and evaluate the gamma-ray intensity from the Galactic Center, instruments with higher angular resolution is needed such as the next generation gamma-ray telescope, Gamma-Ray Large Area Space Telescope (GLAST). N
4 Summary We investigated the dependence of the gamma-ray spectrum from an ADAF on the black hole rotation and examine whether the gamma-ray spectrum can be a probe to investigate the spin parameter. We also applied our model to Sgr A*. We can summarize our results as follows. If the proton gas has a thermal distribution of proton energies, changes in the spin parameter from -0.95 to 0.95 enhance the gamma-ray intensity by orders of magnitude. Therefore in the thermal case, we can estimate the spin parameter from the gamma-ray spectrum using the multi-wavelength observations. If the proton gas forms a power-law distribution of proton energies, the gamma-ray spectrum is not sensitive to changes in the spin parameter, and thus it is not easy to estimate the spin parameter from the gamma-ray spectrum. We applied our model to the Galactic Center. We found that the expected gamma-ray intensities are much lower than the observed value and thus could not find the value of spin parameter. Our model suggests that most of the gamma-ray flux observed by EGRET does not originate in Sgr A*. Acknowledgements K.O. would Like to thank the organizing committees of Galactic Center 2002 and Kobe University for their financial supports.
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References Abramowicz, M. A., Chen, X., Kato, S., Lasota, J.-P., Regev, 0. 1995, ApJ, 438, L37 Baganoff, F. K., et al. 2001, ApJ, submitted (astro-ph 0102151) Bisnovatyi-Kogan, G. S., Lovelace, R. V. E. 1997, ApJ, 486, L43 Dermer, C. D. 1986, ApJ, 307,47 Hartman, R. C., et al. 1999, ApJS, 123,79 Ichimaru, S. 1977, ApJ, 214, 840 Mahadevan, R. 1999, MNRAS, 304,501 Mahadevan, R., Narayan, R., Krolik, L. 1997, ApJ, 486, 268 Mahadevan, R., Quataert, E. 1997, ApJ, 490,605 Manmoto, T., 2000, ApJ, 534, 734 Manmoto, T., Mneshige, S., Kusunose, M. 1997, ApJ, 489, 791 Markoff, S., Melia, F., Sarcevic, I. 1999, ApJ., 522, 870 Melia, F., Fatuzzo, M., Yusef-Zadeh, F., Markoff, S. 1998a, 508, L65 Melia, F., Yusef-Zadeh, F., Fatuzzo, M. 1998b, ApJ, 508. 676 Narayan, R. 2002, in Lighthouses of the Universe, eds. Gilfanov, Sunyaev et al. Springer-Verlag, p 405 Narayan, R., Mahadevan, R., Grindlay, J. E., Popham, R. G., Gammie, C. 1998, ApJ, 492,554 Narayan, R., Yi, I. 1994, ApJ, 428, L13 Narayan, R., Yi, 1. 1995a. ApJ, 444,231 Narayan, R., Yi, I. 1995b, ApJ, 452, 710 Narayan, R., Yi, I., Mahadevan, R. 1995, Nature, 374,623 Oka, K., Manmoto, T. 2003, MNRAS, 340,543 Pohl, M. 1997, A&A, 317,441
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Astron. Nachr./AN 324, No. S1.497-504 (2003) / DO1 10.1002/asna.200385047
The Discovery of Sgr A* W.M. Goss*I, Robert L. Brown'g2,and K.Y. Lo'
' National Radio Astronomy Observatory
' National Astronomy and Ionospheric Center Key words Ga1axy:center-galaxies:individual (Sagittarius)-techniques: interferometer The galactic center compact radio source Sgr A* was discovered on 13 and 15 February 1974 by Bruce Balick and Robert L. Brown using the Green Bank 35 km radio link interferometer (Balick & Brown 1974). We discuss other observationsof this source in the years 1965-1985. Early VLBI observations are described. The name Sgr A* was first used by Robert L. Brown (1982) and has become the accepted name for the compact source at the center of the Milky Way.
1 Introduction The discovery of Sagittarius A as a radio source coincident with the center of the galaxy has been discussed by Goss & McGee (1996). Almost 20 years after the recognition that the center of the galaxy could be associated with Sgr A by Piddington & Minnett (1951) and McGee & Bolton (1954) , Sgr A* was discovered in February 1974 by Bruce Balick and Robert L. Brown in Green Bank, West Virginia. This discovery is certainly one of the more important galactic radio astronomy discoveries of the 1970's and has had wide ramifications during the last 30 years. As an example, the recognition that the radio source Sgr A* is the dim radio source associated with a 2.6 x lo6 Ma black hole has represented a fundamental advance in our understanding of the nuclei of galaxies. The participants in the complex story of the discovery of Sgr A* are numerous: B. Clark, D. Hogg, G. Miley, B. Turner, C. Heiles, R. Ekers, D. Lynden-Bell, D. Downes, A. Martin, B. Balick, R. Brown, M. Goss, K.Y. Lo, U. Schwarz, D. Rogstad and others. (Most of these are still active astronomers.) We present a short history of the discovery process (section 2 ) and provide some details on the naming of Sgr A* in section 3. In section 4, we provide a short summary of the determinations of the secular parallax of Sgr A*. Goss (2003) has also presented a summary of the discovery of Sgr A*. The process of discovery of Sgr A* was the result of the application of a "matched filter" in angular resolution to the properties of Sgr A"; of course, the construction of this filter can only be understood with an a posteriori knowledge of the properties of Sgr A*.
2 The years 1965-1985 In 1966, Clark & Hogg ( 1 966) used the newly completed 2 element Green Bank interferometer at 1 1 cm to investigate the small scale structure of a number of radio sources at I1 cm with a resolution of 10". The source Sgr A was found to have a compact feature with a flux density of 0.3 f.u. (Jy); with this resolution the confusion from Sgr A West is a dominant effect. But these observers were "close" to the discovery of Sgr A". We now know that a resolution of 3" is required at this frequency.
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@ 2003 WILEY-VCH Verlag GmbH & Co KGaA. Wcinheim
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The key observation that led to the discovery of the compact source at the center of the galaxy was an observation by Miley, Turner, Balick, & Heiles (1970). With a baseline of 35 km at 11 cm, these authors discovered a compact source in the H 11 region W5 1 with a Tb > lo5 K. The 42 foot telescope , located at the time of the Miley et al. observations at Huntersville, West Virginia, is shown in Fig. 1. Note the limited tracking capability of this antenna. Although this brightness component in W51 has not been confirmed, the result did set off a string of circumstances that led to the discovery of Sgr A* only four years later. The theoretical framework for the search for evidence of the presence of a compact object at the galactic center was provided by Lynden-Bell (1969) and Lynden-Bell & Rees (1971) , who made the analogy between quasars and the high energy phenomena at the center of the Milky Way: the latter authors propose four tests for the possible detection of a massive object at the center of the galaxy. The second test is :“Very ”. If so Long baseline interferometry may soon be possible with ....as weak as 0.5 f.u. to diameters of it may be possible to determine the size of any central black hole that there may be in our galaxy. However H II may render the central source opaque with a greater angular size.” In the course of 1970 (January and May), Ekers & Lynden-Bell(197 1)used the newly constructed 40 m antenna at the Owens Valley Radio Observatory (Caltech) along with one of the original 90 foot antennas to look for the signature of a black hole in the galactic center. At 6 cm the resolution was 6” by 18”. Ekers & Lynden-Bell detected fine scale structure in the Sgr A West H 11 region. “Although stimulated by the black hole idea our observations are thus more simply explained in terms of young stars and giant H 11 regions.” Again if the resolution had been a factor of about two more favorable , Sgr A* would have been detected. (Goss has pointed out in several lectures in Australia that Ron Ekers has mispelled his name in the acknowledgements to the Ekers & Lynden-Bell paper!) Ekers & Lynden-Bell also performed one of the first interferometer searches for radio recombination lines; they used the 90 foot antenna interferometer at 6 cm. They searched for broad recombination lines from Sgr A (Ekers, private communication); the negative result was not reported in the Ekers & Lynden-Bell paper. This test had also been suggested by Lynden-Bell & Rees (see above) to search for the existence of a massive object at the galactic center. The next step in the quest for compact sources at the galactic center was the result of investigations by Downes & Martin (1971) using the Cambridge One Mile Telescope at 11 and 6 cm with resolutions in RA of 11” and 6”. They describe the overall one dimensional structure (SgrA West and East) with a determination of the spectral indices of the various components. They mention the presence of structures < 10‘‘ in size with flux densities < 1 Jy. Again the discovery of Sgr A* was just over the “resolution horizon”. The discovery of Sgr A* did occur on 13 and 15 February 1974 by Bruce Balick and one of us ( R.L. Brown) using the Green Bank interferometer with an 45 foot antenna at the Huntersville West Virginia site at a distance of about 35 km. This site was the same as the one used in the earlier Miley et al. (1970) observations of W5 1but an improved antenna was now used . This antenna at the Huntersville site is shown in Fig.2; the antenna had a wider range of sky coverage and was operated at the dual frequencies of 11 and 3.7 cm. This interferometer was constructed to serve as a prototype of the Very Large Array which was under construction at the time. The publication of “Intense Sub-arcsecond Structure in the Galactic Center” was published in December, 1974 (Balick &Brown 1974). The resolution at 11 and 3.7 cm was 0.7” and 0.3”, respectively. With this resolution and uv coverage (the three simultaneous baselines from the Green Bank 3 element interferomter and the single antenna at Huntersvile), the extended confusion from the Sgr A West (flux density of 25Jy and size 4U‘)complex was resolved out. Balick & Brown write: “The unusual structure of the sub-arcsecond structure and its positional coincidence with the inner 1-pc core of the galactic nucleus strongly suggests that this structure is physically associated with the galactic center (in fact, defines the galactic center).” These authors compare the compact source with energetic nuclei of other galaxies and even suggest that variations in the radio flux density might be observed. Bob Brown has unearthed a number of fascinating letters from Bruce Balick written during the analysis period from mid March 1974 to 2 May 1974. No copies of Bob’s letters to Bruce were saved. There are no dates on Bruce’s letters. In these letters, Bruce gives in some detail his analysis of the data and their possible interpretations. A number of possible models for the observed visibilities are proposed. Reading N
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these letters today, we are impressed with the meticulous attention to detail in the interpretation of this difficult observation. Here are a few amusing quotes from the letters with an emphasis on a fear (in retrospect somewhat unfounded) of competition. The time frame is toward the end of the period March-May 1974. Bruce writes: “Here are a few thoughts on the 45 foot Sgr A observations. Fred Lo re-analysed some of his VLB observations of Sgr A based on a new position I gave him and found 0.3 f.u. at X 6 cm [see later for a description of this October 1973 observation. The detection was only 2-3 (TI. I think his baseline was Green Bank-Haystack [in fact, Maryland Point]. We’d better publish fast if we want to heat him into print. I haven’t heard from Goss or Downes. Could you call Dave Hogg [then Green Bank site director] and ask if he’s heard anything?”. The following letter expresses some aprehension about IR competetion: “Dave Rank and I [at the University of California at Santa Cruz] are going to try to detect these sources in the IR. Please keep the positions kind of quiet, cause Becklin and co. can wipe us out if they want to. So can Rieke. Your faithful collaborator ...” In 1975, Ekers, Goss, Schwarz, Downes, & Rogstad (1975) combined Westerbork (WSRT) data with Owens Valley Radio Observatory data at 6 cm and made a 2-dimensional image of Sgr A with a resolution of 6“ x 18”. The image was called the WORST image - the Westerbork Owens Valley Radio Synthesis Telescope. In fact the beam shape looked like a sausage - “worst” in Dutch. Sgr A” was just visible at the longest spacings of the interferometer and Sgr A East, the non-thermal source that may be a luminous supernova remnant, was clearly detected as well as hints of the mini-spiral structure of Sgr A West ,an H 11 region associated with the center of the Milky Way. The first VLBI attempt to detect a compact source at the center of the Milky Way was carried out by K.Y. Lo and collaborators in October, 1973. The observation is described by Lo (1974) in his MIT Phd thesis of August 1974: “Interstellar Microwave Radiation and Early Stellar Evolution”. Lo was following up on the Miley et al. W5 1 observation of 1970, to try to confirm the detection of compact structure in HI1 regions, beyond what had been expected theoretically. The observations were at 6 cm between the Green Bank 140 foot (see above) and the Naval Research Labortory 85 foot Maryland Point radio telescope. The GB-MP baseline is mainly E-W with a length of 228 km. A source with a diameter less than 26 mas (EW) 0.1 Jy. As we now know, the size of would be unresolved and detectable down to a flux density of Sgr A* at 6 cm is broadened by interstellar scattering to an EW by NS size of 51 x 27 mas (Lo et al. 1998; Davies, Walsh & Booth 1974). The orientation of the baseline was therefore quite unfavorable for a detection. So, while there were hints of a signal in the visibility amplitude, the detection was not definitive. If the baseline had been oriented in a roughly N-S direction, the source would have been detected. For one of us (Goss), an amusing and somewhat embarrassing episode occurred in the years 1972- 1974. On 2 June 1972, D.Downes and Goss (both working at the Max Planck Institute for Radio Astronomy in Bonn, Germany) submitted proposal D43 to NRAO for an observation of the galactic center with the Green Bank 35 km radio link interferometer. The propsosal was sent to D.Heeschen, D.Hogg and W.E.Howard. We have been able to reconstruct all these events based on the extensive paper archive preserved (in 2003) by Dennis Downes from these pre-email days. The proposal included two positions in Sgr A and three in Sgr B2. A few key sentences from the proposal follow: “In view of the increasing interest in highly collapsed nuclear objects as probable sources of the energy in QSO’s and radio galaxies, it is of paramount importance to pursue investigations of compact structure in Sgr A. Although the center of the galaxy is relatively quiescent, it is so close that we can observe details on a much finer linear scale than is possible in external galaxies, even with VLB techniques. .... We regard this project as an experiment which may he a useful guide to future observations by the VLA These observations might be used in future programs to investigate short-term variability in the galactic center.” Although these ideas were relevant, Goss and Downes were not able to come to the US in 1973-1974 due to problems obtaining travel funds. Also initial observations with the 45 foot telescope were somewhat delayed from late 1972 to mid 1973. Early in 1973, Goss had moved to the Netherlands in a visiting position at the University of Groningen to work on WSRT projects. In addition, Downes was quite busy with early observations with the 100 m Effelsberg N
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telescope. With these pressures, the urgency to complete the Downes-Goss proposal with the Green Bank interferometer decreased. D. Hogg had been in continual contact with Downes and Goss about scheduling. As shown above, Balick and Brown had an earlier NRAO proposal to observe small scale structure (W51 type components) in H 11 regions and Sgr A and Sgr B2 were included. Dave Hogg became aware of the proposal conflict in early 1974 and wrote Downes a letter on 15 February 1974 (note the precise discovery date) proposmg several ways to resolve this conflict. However, Goss and Downes seemed to have lost interest at this point. Of course, the significant result is that Balick and Brown did discover Sgr A* in early 1974 - in fact on 13 and 15 February. The first successful VLBI detection of Sgr A* was made the following year (19 May 1975) by Lo et al. (1975) using the OVRO 40 m and the NASA Goldstone 64-m Mars antenna at 3.7 cm. K.Y. Lo had become a postdoctoral fellow at OVRO after his MIT thesis, but he was interested in following up on the tantalizing hints of detection of Sgr A* on the GB-MP experiment at 6 cm in 1973. It is interesting to recall that after some persuasion by Lo, the observation of Sgr A* was added to the program of his colleagues R. Schillizzi and M. Cohen to study compact symmetric double radio sources. From the California baseline, the inferred size was 20 mas. At an URSI meeting in Boulder, probably in January 1976, after Lo had reported the detection of Sgr Ah at 3.7cm on the OVRO-Mars baseline, Don Backer asked the interesting probing question of how one can be sure that Sgr A* was not a background compact radio source. Interestingly enough, as indicated below, Don Backer answered his own question some years later when he and Dick Sramek detected the secular parallax of Sgr A* due to the rotation of the Sun about the Galactic Center. In the period 10 June 1974 to 10 September 1975, Sgr A * was observed with the early MERLIN array at 0.408,0.96 and 1.66 GHz with baselines of 24 and 127 km. The detections at the latter two frequencies suggested the angular size scales as A’, originating in a turbulent electron distribution along the line of sight (Davies, Walsh & Booth 1974). A number of groups worked on the subsequent VLBI observations of Sgr A* (Kellermann et al. 1977, Lo et al. 1977, Lo et al. 1981, Lo et al. 1985, & Lo et al. 1993). In the 1985 publication of Lo et al., this group determined for the first time that the scattering size of Sgr A* at 3.6 cm was asymmetrical with an axial ratio of 0.55, and at 1.35 cm the limit to the angular size was 20 AU or 2.5 mas. The Green Bank 35 km interferometer was used to determine that the radio spectrum of Sgr A” (Brown, Lo & Johnston 1978) is inverted. Brown & Lo (1982) carried out a ground breaking investigation of the variability of Sgr A* at 11 and 3.7 cm over a time interval of 3 years with 25 epochs (Brown & Lo 1982); time variations were detected over all time scales from days to years. This ground breaking project became the basis for future detailed VLA studies of the various scales in the time variations of Sgr A* (Zhao et al. 2001).
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3 The naming of Sgr A* As far as we can ascertain, the only credit attributed to the naming of Sgr A’ by Brown (1982) is in the Annual Reviews article by Melia & Falcke (2001). The first attempt at a convenient name of the galactic center compact source is by Reynolds & McKee (1980) in a paper entitled : “The Compact Radio Source at the Galactic Center”. This publication presents a model of relativistic outflows, with either a spherical or jet geometry. The fact that the luminosity is 100 times greater than a pulsar but much less than other galactic nuclei was a puzzle. Reynolds & McKee suggest the name GCCRS- the galactic center compact radio source. This name has not survived. Brown & Lo (1982) discuss the variability of Sgr A* (see above) : ”Throughout this paper we use the name Sgr A to refer only and specially to the compact radio source. When necessary, we distinguish this from the more extended radio structure at the galactic center.” In 1982, Backer & Sramek (1982) presented the initial results of the secular motions of Sgr A* using the Green Bank radio link interferometer. The motions were found to be consistent with an object at rest in
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the center of the Milky Way. Backer & Sramek propose the name: SgrA(cn) from “compact non-thermal” object in the galactic center. Again this name has had no staying power. Eight years after the discovery, one of us (Brown) invented the name Sgr A* to distinguish the compact source from the other components in the galactic center and to emphasize the unique nature of this source. Brown (1982) proposed a model of Sgr A* consisting of twin precessingjets with a period of 2300 years. The model has not stood the test of time but the name immediately was accepted. As an example, the VLBI results discussed by Lo et al. (1985) uses the name Sgr A*; the review article by Lo (1987) also uses this nomenclature. Bob Brown provides the following rationale for the name: “Scratching on a yellow pad one morning I tried a lot of possible names. When I began thinking of the radio source as the ‘exciting source’ for the cluster of H IT regions seen in the VLA maps, the name Sgr A* occurred to me by analogy brought to mind by my Phd dissertation, which is in atomic physics and where the nomenclature for excited state atoms is He*, or Fe”, etc.”
4
The motions of Sgr A*
The physical association of Sgr A* with the mass centroid of the nuclear region of the Milky Way remained circumstantial until the observation of the secular motion of Sgr A’ by Backer & Sramek, noted above, was made in 1982 with the Green Bank radio link interferomter. Even in 1983, Martin Rees wrote to Robert L. Brown to report that at the IAU symposium held in June 1983 in Groningen (Netherlands, IAU Symposium No 106) that Jan Oort was worried that the lack of formaldehyde absorption toward Sgr A* could have implied that the radio source is located nearer than the true galactic center. Such concerns were again soundly put to rest based on two companion papers, which were published in the Astrophysical Journal issue of 20 October 1999 (Backer & Sramek 1999; Reid et al. 1999), that summarize the VLA and VLBA determinations of the motions of Sgr A*. Don Backer has pointed out to us that the initial Green Bank 35 km radio link interferometer observation of the 1970’s was inspired by a lunch time conversation with Rick Fisher in about 1975 (see the acknowledgement in Backer & Sramek 1999). The secular parallax due to the motion of the Sun around the center of the galaxy, of course, establishes that Sgr A* is in the galactic center, but more importantly can be used to set a lower limit on the mass of the black hole of a few thousand Ma. In addition, a number of constraints on galactic rotation constants can be determined. The long range goal of the VLBA program (Reid et a1 1999) is the determination of a parallax distance to the galactic center.
5
Summary
The observations of the last 30 years have provided a wealth of information about the source Sgr A* and the environs of the black hole at the center of the Milky Way. Many puzzles remain. We can only imagine the contents of a possible conference on the center of the galaxy that might be held at the time of the 60th celebration of the discovery of Sgr A* in 2034. In March 2004, The National Radio Astronomy Observatory will host a conference in Green Bank, West Virginia, to honor the discovery of Sgr A* exactly 30 years previously and to discuss recent results on this fascinating radio and X-ray source. Acknowledgements The National Radio Astronomy Observatory is a facility of the National Science Foundation operated by Associted Universities, Inc. under cooperative agreement. We thank Dennis Downes, Bruce Balick, Dave Hogg, Don Backer, Ron Ekers and Barry Clark for helpful comments. Much of the work on Sgr A* by Lo was done before his joining the NRAO; he appreciates the important support provided at the critical times in the early years by B. Burke, K. Johnston, J. Moran, A. Rogers, M. Cohen, R. Schillizzi, D. Backer and J. Welch. We thank Pat Smiley for assistance with the Green Bank interferometer photo archive. M. Goss thanks Tony Beasley for his hospilality at the Owens Valley Radio Observatory during February, 2003, while this paper was being written. The late J.H. Oort
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engendered the enthusiasm that initiated the fascination of Goss, Ekers and Schwarz in the 1970's for galactic center research.
References Backer,D.C., & Sramek,R.A. 1982, ApJ, 260,512 Backer,D.C., & Sramek,R.A. 1999, ApJ, 524,805 Basart,J.P., Clark,B.G., &Kramer,J.S. 1968, PASP,80,273 Basart,J.P., Miley,G.K., & C1ark.B.G. 1970, Proc.EEE Trans. Antennas & Propogation, AP-18,375 Balick,B. & Brown,R.L. 1974, ApJ, 194,265 Brown,R.L. 1982, ApJ, 262, 110 Brown,R.L., Lo,K.Y., & Johnston, K.J 1978,AJ, 83, 1594 Brown,R.L., & L0,K.Y. 1982, ApJ, 253, 108 C1arkB.G. & Hogg,D.E. 1966, ApJ, 145,21 Downes,D., & Martin,A.H.M. 1971, Nature, 233,112 Davies,R.D., Walsh,D., &Booth,R.S. 1976, MNRAS, 177,319 Ekers,R.D., Lynden-Bel1,D. 1971, Astrophys. Let., 9, 189 Ekers, R.D., Goss,W.M., Schwarz,U.J., Downes,D., & Rogstad,D.H. 1975, A&A, 43, 159 Fomalont,E.B.2000, in National Radio Astronomy Observatory Workshop Number 27, Radio Interferometry: The Saga and the Science, Proceedings of a Symposium to honor Bany Clark at 60, ed. D. Finley & W.M. Goss (N-kA0),41 Goss,W.M. 2003, in ASP Conf.Ser Vol, Radio Astronomy at the Fringe, ed. J.A. Zensus, M.H.Cohen, & E.Ros (San Francisco: ASP) Goss,W.M., & McGee,R.X. 1996, in ASP Conf. Ser. Vol. 102,The Galactic Center, 4th ESOlCTIO Workshop, ed.R. Gredel (SanFrancisco; ASP), 369 Kellermann,K.I., Shaffer,D.B., Clark,B.G., & Geldzahler,B.J. 1977, ApJ, 214, L61 Lynden-Bel1,D. 1969, Nature, 223,690 Lynden-Bell,D., & Rees,M.J. 1971, MNRAS, 152,461 Lo.K.Y., 1974, Phd thesis, MIT L0,K.Y. 1987, in AIP Conf. Proc No. 155, The Galactic Center, Proc. of the symposiumhonoring C.H. Townes,ed. D.C. Backer (New York AIP), 30 Lo,K.Y., Schilizzi,R.T., Cohen,M.H., & Ross,H.N. 1975, ApJ, 202, L63 Lo.K.Y., Cohen,M.H., Schilizzi,R.T.,& Ross,H.N. 1977, ApJ, 218, 668 Lo,K.Y., Cohen,M.H., Readhead, A.S.C., & Backer, D.C. 1981,249,504 Lo,K.Y., Backer,D.C., Ekers,R.D., Kellermann,K.I., Reid,M.,& Moran, J.M. 1985, Nature, 315, 124 Lo,K.Y., Backer,D.C., Kellermann,K.I., Reid,M., Zhao,J.-H., Goss, W.M., &Moran,J.M. 1993, Nature, 362, 38 Lo.K.Y., Shen,Z.-Q., Zhao,J.-H., & H0,P.T.P. 1998, ApJ, 508, L61 Melia,F. & Falcke,H. 2001, ARA&A, 39,309 McGee,R.X., Bo1ton.J.G. 1954, Nature, 173,985 Miley,G.K., Tumer,B.E., Balick,B., & Heiles,C. 1970,ApJ. 160, L119 Piddington,J.H. & Minnett,H.C. 1951, Aus J Sc. Res., A4,495 Reynolds,S.P. & McKee,C.F. 1980, ApJ, 239,893 Reid,M.J.,Readhead,A.C.S., Vermeulen,R.C., &Treuhaft,R.N. 1999, ApJ, 524,816 Zhao,J.-H., Bower, G.C. & Goss, W.M. 2001, A@, 547, L29
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Fig. 1 The 42 Foot antenna. The initial location of this antenna was at Spencer's Ridge - 11 km NE of the Green Bank interferometer. The operation was at 1 1 cm. Basart et al. 1968 and Basart et al. 1970 describe initial observations with a two element interferometer consisting of this antenna combined with one of the 85 Foot antennas at Green Bank. The sky coverage was limited to declinations from 0 'to +66 'and hour angles within 2h40"'n of the meridian. For the observations of W51 described by Miley et al. 1970, the antenna had moved to the Huntersville site - 35 km to the SW of Green Bank. These observations were a partial inspiration for the Sgr A* observations of 1974. Later the 42 Foot antenna was replaced by the fully steerable antenna shown in Fig.2 at the Huntersville site. Fomalont (2000) summarizes the development of the Green Bank radio link interferometer in the years 1966 to 1978.
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Fig. 2 The 45 Foot telescope used at Huntersville, West Virginia as the 3.5 km outstation for the NRAO Green Bank radio link interferometer. The discovery of Sgr A* was made in February, I974 by Bruce Balick and Robert L. Brown using this instrument (Balick & Brown 1974). The dual frequency instrument operated at 11 and 3.7 cm and was used as a prototype for the VLA in the planning stages during the early 1970’s (see Fomalont 2000). The 45 Foot telescope was the fourth element of the interferometer ; correlations were performed with all three of the 8.5 Foot antennas at Green Bank. This smaller antenna is now at the Green Bank site and has had an illustrious career as a component of the tracking stations for the HALCA VLBI spacecraft. The two antennas for telemetry are pointing to the Green Bank site to the northeast; the local oscillator signal was transmitted by a two way link at 1.3GHz while the 18GNz link was used for telemetry and IF transmission. During the summer there was no clear line of sight to the main interferometer site due to leaves in intervening trees. A passive reflector was used on a hill behind the main site to overcome this problem. Much of the development work for the radio link was done by N.G.V. Sarma who was on a sabbatical from the Tata Institute for Fundamental Research (Ooty) in In&a
Astron. Nachr./AN 324, No. S1.505 -51 1 (2003)/ DO1 10.1002/asna.200385084
The Position, Motion, and Mass of Sgr A* Mark J. Reid * I , Karl M. Menten’, Reinhard Genze13, Thomas Ott3, Rainer Schode13, and Andreas Brunthaler’ ’ Harvard-Smithsonian Center for Astrophysics, 60 Garden St., Cambridge, MA 02138, U.S.A. Max-Planck-Insititut fur Radioastronomie,Auf dem Hugel 69, D-53121 Bonn, Germany Max-Planck-Insititut fur extraterrestriche Physik, Giessenbachstrasse,D85748 Garching, Germany
Key words Sgr A*, black holes, proper motions, SiO masers
Abstract. We report progress on measuring the position of Sgr A* on infrared images, placing limits on the motion of the central star cluster relative to Sgr A*, and measuring the proper motion of Sgr A* itself. The position of Sgr A* has been determined to within 10 mas on infrared images. To this accuracy, the gravitational source (sensed by stellar orbits) and the radiative source (Sgr A*) are coincident. Proper motions of four stars measured both in the infrared and radio indicate that the central star cluster moves with Sgr A* to within 70 km sC1 . Finally, combining stellar orbital information with an upper limit of 8 km s-’ for the intrinsic proper motion of Sgr A* (perpendicular to the Galactic plane), we place a lower limit on the mass of Sgr A* of 4 x lo5 Mo.
1 Introduction The precise position and proper motion of Sgr A* are of fundamental importance in order to understand the nature of the super-massive black hole (SMBH) candidate and its environment. Unfortunately, Sgr A* lies behind about 30 mag of visual extinction, and currently it can only be detected in the radio, infrared, and x-ray bands. While its radio emission is easily detected, the same cannot be said for its infrared and x-ray emission. In both of these wavebands, emissions from nearby (in angle) stars make it difficult to 10 milli-arcseconds isolate and measure the emission from Sgr A*. Only with positions accurate to (mas) can one confidently separate Sgr A* from confusing stellar sources and determine its spectral energy distribution and time variations. Stellar proper motions, accelerations, and even orbits are now being determined to high accuracy at infrared wavelengths, and the position of the central gravitational source (presumably Sgr A*) can be measured to mas accuracy. If Sgr A* is indeed a SMBH, then the gravitational source, inferred from 10 stellar orbits, and the radiative source, directly seen in the radio band, should coincide to within Schwarzschild radii (10R,,h E 0.08 mas M 1013 cm for Sgr A*). Thus measuring the position of Sgr A* in the infrared to sub-mas levels is of fundamental importance in testing the SMBH paradigm. The apparent proper motion of Sgr A* directly determines the sum of the angular rotation speed of the and any peculiar motion of Sgr A* ( V s g T ~with * ) respect Sun about the Galactic center, ( 0 0 VO)/Ro, to the dynamical center of the Galaxy. Thus, Sgr A*’s proper motion can provide a direct measurement of Galactic rotation. In addition, the combination of stellar motions and an upper limit on the motion of Sgr A* itself can yield a strong lower limit to the mass of the SMBH candidate. This paper reports recent progress on locating Sgr A* on infrared images, measuring its proper motion, and placing a lower limit on its mass. N
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* Corresponding author: e-mail.
[email protected],Phone: +1617495 7470, Fax: + I 617495 7345 0 ZW3WILEY-VCH Verlag GmbH & Co
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Fig. 1 Location of Sgr A* on a July 1995 2pm wavelength image of the inner 2 arcsec of the Galactic center, adapted from Menten et al. (1997) by Reid et al. (2003). The circle centered at the position of Sgr A* has a radius of 15 mas, corresponding to a la position uncertainty.
2 Previous Results Menten et al. (1997) detected SiO maser emission from red giant and supergiant stars within 12 arcsec of Sgr A*. Since the maser emissions originate from within about five stellar radii of the host star, they can serve to precisely locate the star at radio wavelengths relative to the strong radio source Sgr A*. Also, the SiO maser stars are very bright at infrared wavelengths, and one can use the radio positions of two or more stars to calibrate the infrared plate scale and rotation and then align the infrared image with the radio image containing Sgr A*. Following this method, Menten et al. located Sgr A* on a 2 pm wavelength image to an accuracy of 30 mas (lo).No source of emission was seen at the position (see Fig. 1) of Sgr A*, and an upper limit of 9 mJy (de-reddened) was established. Reid et al. (1999) and Backer & Sramek (1999) published observations of the apparent proper motion of Sgr A*. Both papers show that Sgr A* appears to move toward the south-west along the Galactic plane at about 6 mas yr-l. This is consistent with the angular rotation rate of the Sun in its 220 Myr period about
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the Galaxy. Removing the effects of the Sun's orbit yields an upper limit to the peculiar motion of Sgr A* of about 20 km s-l . Reid et al. interpreted this upper limit to indicate that the mass of Sgr A* exceeds lo3 Ma, ruling out any stellar source.
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3 Recent Advances 3.1 The Infrared Reference System Recently, a wide-format infrared camera (CONICA; Lenzen et al. 1998) with an adaptive optics assisted imager (NAOS; Rousset et al. 2000) was installed on one of the ESO 8.2-m VLT telescopes. This has produced excellent data for diffraction-limited imaging. Very deep images of the Galactic center with a field of view of 28 arcsec were taken with these instruments at 2 pm wavelength early in 2002. These images proved to be of excellent quality. Over the past seven years, radio frequency observations of SiO maser sources in the Galactic center have been conducted with the NRAO VLBA and VLA. Both telescopes have been used to measure positions and proper motions of SiO maser stars with accuracies of about 1 mas and 1 mas yr-', respectively, relative to Sgr A*. Seven maser stars within 15 arcsec of Sgr A* have now been measured to these accuracies (Reid et al. 2003). By combining the radio positions of seven maser stars with their apparent positions on the new VLT images, we could determine the infrared plate scales and rotations with high accuracy. After aligning the radio and corrected infrared images, we found the residual differences in the maser star positions to be about 6 mas. This verified that the new infrared images have very small distortions across the field of view and allowed us to determine the position of Sgr A* with a 10 mas (la)uncertainty (Reid et al. ). Figure 2 shows a 2 p m wavelength image taken on 2 May 2002, with the position of Sgr A* and two nearby stars indicated. At this time, the fast moving star S2 was near pericenter in its orbit about Sgr A* (Schodel et al. 2002). The proximity of S2 to Sgr A* at this time (16 mas) precludes any significant measurement of the flux density of Sgr A*. The location for Sgr A* is within 10 mas of the gravitational source inferred from orbital solutions for the star S2 (Schodel et al. ; Ghez et al. 2003). Thus, the radiative source (Sgr A*) and the gravitational source of the SMBH candidate are co-located to within about 1000 Rsrh. Previous infrared proper motions have been relutive motions, with the motion reference defined by setting the average of large numbers of stellar motions to zero. We have compared the radio and infrared proper motions directly in order to transfer the infrared motions to a reference frame tied to Sgr A* (Reid et al. 2003). In principle, one can make this reference frame transfer using a single star with well determined motions. However, we chose to average the results from the four SiO maser stars within 10 arcsec of Sgr A* that have measured proper motions both in the radio and infrared. The unweighted mean difference (and standard error of the mean) of these stars is 0.8410.85 mas yr-' toward the east and -0.2510.96 mas yr-l toward the north. Since I mas yr-' corresponds approximately to 38 km s-I (for a distance to the Galactic center of 8.0 kpc), we conclude that the central star cluster moves with Sgr A* to within about 40 km sP1 per coordinate axis, or within about 70 km s-l for a 3-dimensional motion. 3.2
Proper Motion of Sgr A*
The apparent proper motion of Sgr A*, with respect to extragalactic radio sources, was measured with the VLBA by Reid et al. (1999). In that program, Sgr A* was used as a phase reference to calibrate the interferometer phases for two compact radio sources. The location of these sources is shown in Figure 3. Relative to an extragalactic source, Sgr A* would be expected to move toward the south-west, mostly along the Galactic plane (as depicted in Fig. 3), owing to the orbit of the Sun about the Galactic center. Figure 4 shows the position residuals of Sgr A* relative to the compact extragalactic source 51745283. The positions for 1995 through 1997 are from Reid et al. (1999). Those for 1998 through 2000 are new measurements. As one can see, the apparent motion of Sgr A* continues along the Galactic plane. The dominant term in the apparent motion of Sgr A* comes from the orbit of the Sun. This can be decomposed
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Fig. 2 Location of Sgr A* on 2 May 2002 on an infrared (2 p) image of the inner 2 arcsec of the Galactic center, adapted from Reid et al. (2003). The circle centered at the position of Sgr A* has a radius of 10 mas, corresponding to a 117position uncertainty. The star S2 was near pencenter in its orbit about Sgr A* when this image was taken (Schodel et al2002; Ghez et al. 2003)
into a circular motion of the local standard of rest, Oo/Ro = 220 k m sC1/8.O kpc (see Reid et al. 1999for details), and the peculiar motion of the Sun, Vo/& M 20 km s-'/8.0 kpc. Removing these terms from the observed proper motion, yields estimates of the peculiar motion of Sgr A*.
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While we currently do not know the component of 00 Va in the plane of the Galaxy to better than about 10 to 20 km s-l , we do know the component perpendicular to the Galactic plane to better than 1 km s-l . Since the circular motion of the LSR is, by definition, entirely in the plane of the Galaxy, the only contribution to the apparent motion of Sgr A* perpendicular to the plane of the Galaxy is the Z-component of the Sun's peculiar motions, V z o . This component can be estimated by averaging the motions of very large numbers of stars in the Solar neighborhood, which should directly indicate -Vzo. An estimate, using the Hipparcos database, indicates Vza = 7.16 i 0.38 km s-l toward the north Galactic pole (Dehnen & Binney 1998). After removing this contribution to the apparent motion of Sgr A*
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Fig. 3 Positions of Sgr A* and two compact extragalactic radio sources superposed on a 90 cm wavelength image of the Galactic center region made with the VLA by LaRosa et al. (2000). The expected motion of Sgr A*, owing to the orbit of the Sun about the Galactic center, is indicated by the arrow.
perpendicular to the plane of the Galaxy, we arrive at an estimate of 5 i3 km s-l for this component of Sgr A*’s peculiar motion. This result significantly improves on the limits given by Reid et al. (1999) and Backer & Sramek (1 999).
4
The Mass of Sgr A*
From infrared observations of stellar orbits (Schodel et al. 2002, Ghez et al. 2003), we know that a mass of M 3 x lo6 Ma is contained within a radius of = 100 AU. With this information, and an upper limit on the Z-component of the velocity of Sgr A*, V,, for which we adopt 8 km s-’ , one can estimate a lower limit to the mass of Sgr A*. The basic parameters of the problem are the total enclosed mass, M p n r ( R )including , a possible SMBH and stars with typical individual mass, m, that are enclosed within a radius, R, and an upper limit on the Zcomponent of the velocity, V,, of a “test” object (Sgr A* in our case) of mass M . In the past, two limiting cases of mass estimators have been discussed for this problem: equipartition of kinetic energy (see Backer
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East Offset (mas) Fig. 4 Position residuals, with lu error bars, of Sgr A* relative to J1745-283 on the plane of the sky. Each measurement is indicated with an ellipse, approximating the apparent scatter broadened size of Sgr A* at 43 GHz. The dashed line is the variance-weighted best-fit proper motion, and the solid line gives the orientation of the Galactic plane. The expected position of Sgr A* at the beginning of each calendar year is indicated.
& Sramek 1999) and momentum (see Reid et al. 1999). Equipartition of kinetic energy implies that
MV2
N
mu2 ,
(1)
where v2 NN GM,,,(R)/R is a characteristic stellar velocity at radius R (which must be great enough so that the mass in stars exceeds that of the the test object, i.e., Menc(R)2 2M). Equipartition of energy is both theoretically and observationally well founded for the case of stellar clusters. However, for the case of a dominant central mass, which could greatly exceed the total mass of stars (within a given radius), Reid et al. argued that one would be dealing with true orbits and that equipartition of momentum would then be appropriate: MVwmu. (2)
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It turns out that both estimators are correct, but for answering different que3tions. If one asks what is the expected velocity of a SMBH that is perturbed by close passages of stars which orbit it, then the momentum equation (2) applies. This is almost surely the case for Sgr A* and nearby stars such as S 1 and S2. For star S2, which has a mass rn of x 15 Mo (Ghez et al. 2003), during pericenter passage t i M 6500 km s-l and Eq. (2) implies that one would expect a 3 x lo6 M a SMBH’s peculiar motion to be V N 0.03 kni s-l . Following this approach one can calculate an extremely conservative lower limit for Sgr A*’s mass. While this is valid, it is not an optimum estimate. Alternatively, if we ask for the minimum mass of a central object which does not totally dominate the within a given radius R, and which complies with the observed velocity limit, enclosed mass, hfenC(R), then we have a different case. For this case, where the enclosed mass in stars within R is comparable to or exceeds the mass of Sgr A*, M , equipartition of kinetic energy should apply. When evaluating Eq. (1) one must use velocities for stars with radii near R. Conceptually, as the velocity limit for Sgr A* improves, the estimated mass limit increases quadratically in V. This continues until the estimated mass dominates over the stellar component and our assumption is violated. At this point, however, one has already ascribed most of the enclosed gravitational mass to Sgr A*. A recent paper by Chatterjee, Hernquist & Loeb (2002) analyzes our mass estimation problem in a manner similar to that described above. They assume a black hole at the center of a stellar cluster, which is distributed in space according to a Plummer profile with a characteristic scale a. The rnininiurn black hole mass occurs for a approximately equal to the radius, R, within which the enclosed mass is measured. In this case, the mass estimator (their equation 42) can be simplified to the following:
provided V,” > Gm/R,which is met for V, = 8 km s-’ , 711 = 1 Ma,and R = 100 AU. Then for the observed M,,,(R) = 3 x lo6 Ma, Eq. 3 gives a lower limit to the mass of Sgr A*of M 2 . 4 x los M a . Our lower limit to the mass of Sgr A* is now within about a factor of 10 of the total mass required by recent IR data. Since the uncertainty in the proper motion (cv)can decrease with the spanned observing ’7 ; (until the limit approaches time ( T )as ov c( T-”’, and the lower limit to the mass of Sgr A* scales as ~ the total enclosed mass), we can expect improvement in the limit for A4 cx T 3 over the next few years. When we reach a motion limit of 1 to 2 km s-l for Sgr A*, then essentially all of the mass sensed gravitationally by stellar orbits must come from Sgr A* itself. Should future VLBI measurements at 5 1inm wavelength show that the intrinsic size of Sgr A* is 5 0.1 AU, then we may be in a position to conclude that for Sgr A* most of the mass required for a SMBH is contained within a few Rsch!
References Backer, D. C. & Sramek, R. A. 1999, ApJ, 524, 805 Chatterjee, P., Hernquist, L., & Loeb, A. 2002, ApJ, 572,37 1 Dehnen, W. & Binney, J. J. 1998, MNRAS, 387 Ghez, A. et al. 2003, to appear in ApJ(Lett.) LaRosa, T. N., Kassim, N. E., Lazio, T. J. W., & Hyman, S. D. 2000, AJ, 119,207 Lenzen, R., Hofmann, R., Bizenberger, P., & Tusche, A. 1998, Proc. SPIE, 3354, 606 Menten, K. M., Reid, M. J., Eckart, A,, & Genzel, R. 1997, ApJ(Lett.),475, L111 Reid, M. J., Readhead, A. C. S., Verrneulen, R. C., & Treuhaft, R. N. 1999, ApJ, 524, 816 Reid, M. J., Menten, K. M., Genzel, R., Ott, T., Schodel, R. & Eckart, A. 2003, to appear in ApJ, 587 Rousset, G. et al. 2000, Proc. SPIE, 4007, 72 Schodel, R. et al. 2002, Nature, 419, 694
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Astron. Nachr./AN 324, No. S1.513-519 (2003)/ DO1 10.l002/asna.200385098
Tidal processes very near the black hole in the Galactic Center Tal Alexander Faculty of Physics, The Weizmann Institute of Science, POB 26, Rehovot 76100, Israel
Key words Galaxy: center, galaxies: nuclei, Black hole, stars: kinematics Abstract. The accessibility of the 3 x loGA40 massive black hole (MBH) in the Galactic Center (GC) offers a unique opportunity to probe a variety of strong tidal interactions of stars with a MBH or with other stars in the high density cusp around a MBH. We show that such interactions can affect a significant fraction of the stellar population within the MBH radius of influence. We consider three processes that could possibly modify stellar structure and evolution there. (1) Tidal spin-up by hyperbolic star-star encounters. (2) Tidal scattering of stars on the MBH. (3) Tidal heating of tidally captured inspiraling stars-"squeezars". We discuss the implications for stellar populations near MBHs and for the growth of MBHs by tidal disruption, and the possible signatures of such processes in the GC. We compare the event rates of prompt tidal N
encounters (tidal disruption and tidal scattering) with slow inspiral events (squeezars / tidal capture), and find that, contrary to what was assumed in past studies, tidal capture in the presence of scattering is an order of magnitude less efficient than prompt disruption and so does not contribute significantly to the growth of the MBH.
1 Introduction Strong tidal interactions involving stars are expected to occur frequently near a MBH in a galactic center. First, the MBH is a mass sink, which drives a flow of stars from the MBH radius of influence T h to the center, to replace those it has destroyed. An inevitable consequence of this flow is that some stars are deflected into orbits whose periapse r p lies just outside the critical radius for destruction. We focus here on MBHs like the one in the GC, whose mass m is small enough so that the event horizon T , lies inside the where M, and R, are the stellar mass and radius 108Mo tidal disruption radius T t R*(T~/M*)'/~, for a solar type star). Such stars will suffer an extreme tidal impulse, but will not be destroyed, at least not on their first peri-passage. There are two possible outcomes: that the star is ultimately disrupted, or that it avoids subsequent encounters with the MBH. Both will be considered here in detail. Second, a variety of formation scenarios predict that MBHs should lie in the center of a high density stellar cusp (e.g. Bahcall & Wolf 1977; Young 1980). The diverging stellar density implies that there must be some volume around the MBH where close tidal encounters occur on timescales significantly shorter than the typical stellar lifetime. Such encounters will have a very different outcome from those that occur in globular clusters that do not contain a MBH. In most cases the encounters will not lead to tidal capture. Instead the two stars will continue on their separate ways after experiencing a brief strong tidal impulse. Extreme tidal interactions, which transfer energy and angular momentum from the orbit to the star, can affect its structure and subsequent evolution by heating it, spinning it up, mixing it, or ejecting some of its mass. This is interesting in view of the observed presence of unusual stellar populations near MBHs: the blue nuclear cluster in the inner 0.02 pc of the GC (Genzel et al. 1997), and around the MBH in M3 1 (Lauer et al. 1998); evidence for anomalously strong rotational dredge-up in an M supergiant near the MBH in the GC (Can; Sellgren & Balachandran 2000), but not in a high density nuclear cluster without a MBH (Ramirez et al. 2000); the unusually high concentration of very rare extreme blue He supergiants around the Galactic MBH (Krabbe et al. 1991; Najarro et al. 1994). The observable signature of an extreme tidal interaction on a star cannot be predicted with certainty at this time, although some reasonable conjectures can be made (Alexander & Livio 2001; Alexander &
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Morris 2003). It is clear however, as shown below, that the tidal energy deposited in the star can reach a significant fraction of its binding energy, and that the angular momentum extracted from the orbit can spin it up to a significant fraction of its break-up velocity. We therefore proceed on the assumption that the effects can be observationally interesting and explore the dynamical processes that give rise to such interactions. Furthermore, tidal disruption and collisional stellar mass loss are important channels for supplying mass to a low-mass MBH (Murphy, Cohn & Durisen 1991), and so extreme tidal processes may provide observable links between the properties of the stellar population near the MBH and its evolutionary history. At a distance of 8 kpc (Reid 1993), the low mass 3 x lo6 Ma MBH in the GC (Ghez et al. 2000; Schodel et al. 2002) is the nearest and observationally most accessible MBH. Although heavily reddened, deep high resolution astrometric, photometric and spectroscopic IR observations of thousands of stars near the MBH provide information on their luminosity, colors and orbits (e.g. Eckart, Ott & Genzel 1999; Figer et al. 2000; Gezari et al. 2002). The GC thus offers a unique opportunity to probe a variety of close tidal interactions of stars with a MBH or with other stars in the high density cusp around it.
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2 Tidal Spin-up by star-star encounters Stars moving rapidly around a MBH in a dense stellar cusp will suffer numerous hyperbolic tidal encounters over their lifetimes. Although such encounters transfer some orbital energy and angular momentum to the colliding stars, they rarely remove enough energy for tidal capture. This is in marked contrast to collisions in high density cores of globular clusters without a MBH, where the colliding stars are on nearly zero-energy orbits and close collisions result in tidal binary formation. The effects of hyperbolic encounters on the stars are mostly transient. The stellar dynamical and thermal timescales are very short compared to the mean time between collisions, and so apart from some mass-loss in very close collisions, the star is largely unaffected. It is however more difficult for the star to shed the excess angular momentum, since magnetic breaking operates on timescales of the order of the stellar lifetime (Gray 1992). High rotation is therefore the longest lasting dynamical after-effect of a close encounter. Over time, the stellar angular momentum will grow in a random walk fashion due to successive, randomly oriented tidal encounters. We consider the tides raised by an impactor star of mass m on a target star of mass M, and radius R, as m follows an unbound orbit with periapse rP from M,. - We will use the tilde symbol to express quantities in - units where G = M* = R, = 1. In these units, 0 = 1 is the centrifugal break-up angular frequency, Eb = -1 is the stellar binding energy, up to a factor of order unity, and Ft = 6 * / is3 the tidal radius. The orbital energy A E and angular momentum AJ”that are transferred from the orbit to the star by an impulsive tidal encounter are related by AE= fipAY, where fip is the relative angular velocity at periapse. Assuming for simplicity rigid body rotation, the change in the stellar angular velocity due to one parabolic encounter is given to leading order in the linear multipole expansion by (Press & Teukolsky 1977)
where i i s the stellar moment of inertia (assumed to remain constant), Tz is the C=2 tidal coupling coefficient (calculated numerically for a given stellar structure model), and q= [ 3a from linear motion). Three of them, S l , S2, and S8 are well known from previous publications (Ghez et al. 2000; Eckart et al. 2002; Schodel et al, 2002). The other three sources, S12, S13, and S14, are fainter ( K 2 15) sources, the proper motions of which could only be disentangled from the high confusion central stellar cluster with a sufficiently large data base such as presented here (see also Schodel et al. 2003).
4 Radial anisotropy We examined the Sgr A* stellar cluster using Y T R = (v? - v g ) / w zas anisotropy estimator, where v is the proper motion velocity of a star, and VT and U R its projected tangential and radial components. A value of +1 signifies projected tangential motion, -1 projected radial motion of a star. The properties of the anisotropy parameter T T R are discussed in detail in Genzel et al. (2000). They show that an intrinsic threedimensional radiavtangential anisotropy is reflected in the properties of the two-dimensional anisotropy estimator Y T R . In Figure 1 we show a histogram of the parameter Y T R for different sub-samples (distance to Sgr A* 5 0.6", 5 LO", and 5 1.2") of our proper motion data for the epoch 2002.7. The projected velocities of the accelerated stars at this epoch were estimated by linear fits to sufficiently short parts of their trajectories. Repeating the anisotropy analysis for the epoch 1995.5 (where some of the stars had significantly different positions and velocities) does not change the distribution of counts in the histrograms significantly. The number of stars on projected radial orbits is 2 - 3a above the number of stars on projected tangential orbits (taking Poisson errors). The number of stars on projected tangential orbits decreases significantly with decreasing distance to Sgr A*. More proper motion data are needed in order to settle the question of anisotropy, but the present analysis presents a very intriguing result. Should the radial anisotropy of the Sgr A* cluster indeed be proven to be true with the larger proper motion samples expected from future observations, theoretical and modeling efforts will be needed to understand this property of the Sgr A* stellar cluster. As a bottom line, we want to point out that the general distribution of the anisotropy parameter definitely excludes a tangentially anisotropic cluster. A significant tangential anisotropy would be expected in systems with a binary black hole, where stars on radial orbits would be ejected or destroyed preferentially (see e.g. Gebhardt et al. 2002).
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Stellar orbits
The NACO GC observations in springhmmer 2002 covered the pericenter passage of the star S2 around Sgr A* in a tightly sampled time series. Combining these observations with SHARP imaging data since 1992 (taken from Ott et al. 2003), Schodel et al. (2002) determined a unique keplerian orbit for S2. In the left panel of Figure 2 we compare the orbit of S2 of Schodel et al. (2002) with the orbit of S2 as determined in the present work. There are three important differences between the two analyses (see Schodel et al. 2003): ( 1 ) Here, the SHARP positions were obtained with different data reduction and analysis techniques (from a comparison with Ott et al. 2003 we estimated an overall systematic error of -3 mas). (2) Schodel et al. (2002) measured the positions of S2 from one final shift-and-add image for each NACO epoch and estimated the errors conservatively. Here, the S2 position for each NACO observing epoch results from measurements on several tens of individual short-exposure NACO images, with the standard deviation taken as error. (3) We treated the projected position of the focus of the elliptical orbit as a free parameter in the fit (see Schodel et al. 2003). The two analyses compare very well, with the determined orbital parameters agreeing within the errors. In our present analysis, we obtain a central mass of 3.3 f 0.7 x 106Ma. The position of the acceleration
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center is offset a mere 2.0 f2.4 mas East and 2.7 k 4.5 mas South of the nominal radio position of Sgr A*, i.e. clearly within the error circle of the radio measurement. This strongly supports the assumption that the dark mass is indeed coincident with Sgr A*. The orbit has an eccentricity of 0.87 f0.02, an inclination of 45.7 f 2.6 degrees, a period of 15.7 f 0.74 years, a semi-major axis of 4.54 f0.27 mpc, and a pericenter distance of 0.59 f0.10 mpc. Significant sections of the orbits were observed as well in the case of the stars S12 (pericenter passage in 1995.3) and S14 (pericenter passage in 1999.9). However, the constraints on their orbits from our data are not very tight. S14 is identical with the source SO-16 of A. Ghez et al. (priv. c o r n . ) , who first determined an orbit for this source. S14 is on an extremely eccentric (e = 0.97 f0.05) and highly inclined (i > SOdeg) orbit and approaches Sgr A* to within -0.4 mpc (S2: 0.6 mpc). In principle, its orbit would allow to constrain the central mass distribution even tighter than S2. Unfortunately, the uncertainty in the orbital parameters of S14 resulting from our data is too high for this purpose. The orbital segments observed for the stars S1 (pericenter passage around 1999/2000), S8, and S13 are too small for determining a unique set of parameters for them, but we constrained them by using fixed values for the inclination angle (see Schodel et al. 2003). Approximate values of the inclination of the orbital planes could be estimated from the measured acceleration of the stars and the well known mass of the central dark object (see Figure 3). We plot all analyzed six orbits in the right hand panel of Figure 2. All analyzed orbits have moderate to high eccentricities. Future measurements of orbital eccentricities of more stars near Sgr A* will allow testing for anisotropy of the central cluster (see Schodel et al. 2003).
6 Nature of the enclosed mass With the measured proper motions within 1.2” of Sgr A* we calculated Leonard-Merritt (LM, Leonard and Merritt 1989) estimates of the enclosed mass. In order to take the strongly variable velocity of the 6 stars with significant acceleration into account, we produced various velocity lists for the analysis, where we estimated the projected velocity of these stars at different epochs. From the diffrent lists, we obtain an average LM mass of 3.4 f 0.5 x 1O6M@(for details see Schodel et al. 2003). This agrees well with the mass estimate from the orbit of S2.
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Fig. 2 Left panel: The orbit of S2 (black) as determined by Schodel et al. (2002) compared with the orbit as determined in the present work (green). The black cross and circle denote the radio position of Sgr A* and its errors. The red cross designates the position of the focus of the orbit and its error resulting from the present analysis. Right panel: Currently, we can determinekonstrain the orbits of 6 stars near Sgr A*: S1, S2, S8, S12, ,513, and S14. all orbits have
moderate to high eccentricities. Figure 3 is a plot of the measured enclosed mass against distance from Sgr A*, in close analogy to Figure 17 of Genzel et al. (2000) and Figure 3 of Schodel et al. (2002). The data show that the central mass distribution is remarkably well described by the potential of a point mass over 3 orders of magnitude in spatial scale, from 0.8 light days to 2 light years. The contribution of the extended stellar cluster around Sgr A* to the total mass cannot be more than mostly a few hundred solar masses within the peri-center distance of S2 (Mouawad et al. 2003). Fitting a model composed of a point mass plus the visible outer stellar cluster with a core radius of 0.34 pc and a power-law slope of a = 1.8 to the data gives a value of 2.9 0.2 x 106Ma for the central dark mass. This agrees within the errors with the LM mass estimate of the innermost stars and with the masses calculated from the orbital parameters of S2 and S12. It is higher than the 2.6 0.2 x 106Ma given by Schodel et al. (2002), but the two values agree within their errors. The main differences of the present analysis to Schodel et al. (2002) are: (1) The error of the mass estimate from the orbit of S2 has been reduced by taking the position of the orbital focus explicitly into account. (2) The innermost LM mass estimate of Schodel et al. (2002) was based on the Ott et al. (2003) data. It has been replaced by the LM mass estimate from the present work, which is based on a more abundant 1" of Sgr A*. (3) The LM mass estimates in Figure 3 of Schodel et al. data base in the region within (2002) and Figure 17 of Genzel et al. (2000) were corrected downward by 510% because they assumed a power-law slope of a = 1.8 for the stellar cluster in the innermost few arcseconds. Here, we use a power-law slope of a M 1.4 for the stellar cusp around Sgr A* (Genzel et al. 2003). This means that the LM mass estimates have previously been underestimated by 10%. The orbit of S2 places very tight constraints on the distribution of the central dark mass: If the central point mass were replaced by a Plummer model cluster of dark astrophysical objects, its central mass density would have to exceed 2 . 2 ~ 1 0 ' ~ M ~ p calmost - ~ , 5 orders of magnitude greater than previous estimates
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(Ghez et al. 1998,2000; Genzel et al. 2000). The lifetime of such a hypothetical cluster would be < 10' years (Maoz 1998). An alternative model to supermassive black holes in galactic nuclei are balls of heavy, degenerate neutrinos (Tsiklauri and Viollier 1998; Munyaneza and Viollier 2002). In order to explain the whole mass range of dark central objects in galaxies with such a model, the neutrino mass cannot be higher than 17keV (Melia and Falcke 2001). However, the orbital parameters of S2 would demand a neutrino mass of > 5OkeV in the case of the dark mass in the Galactic Center. The only dark particle matter explanation that cannot b e ruled out by the present data is a ball of bosons (Torres et al. 2000). However, it would be hard to understand how the bosons first manage to reach such a high concentration, and then avoid forming a black hole by accretion of the abundant gas and dust in the GC. We therefore conclude that the most probable form of the dark mass at the center of the Milky Way is a single, supermassive black hole.
Acknowledgements We like to thank the ESO NTT team for their help and support during ten years of observations with the SHARP guest instrument. We thank the NAOS and CONICA team members for their hard work, as well as the staff of El Paranal and the Garching Data Management Division for their support of the commissioning and science verification. Based on observations obtained at the Gemini Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under a cooperative agreement with the NSF on behalf of the Gemini partnership: the National Science Foundation (United States), the Particle Physics and Astronomy Research Council (United Kingdom), the National Research Council (Canada), CONICYT (Chile), the Australian Research Council (Australia),CNPq (Brazil) and CONICET (Argentina).
References Chakrabarty, D. and Saha, P. 2001, AJ, 122,232 Eckart, A., Genzel, R., Krabbe, A., Hofmann, R., van der Werf, P.P., Drapatz, S. 1992, Nature, 355,526 Eckart, A,, Genzel, R., Hofmann, R., Sams, B.J., and Tacconi-Garman L.E. 1995, ApJ, 445, L23 Eckart, A. and Genzel, R. 1997, MNRAS, 284,576 Eckart, A., Genzel, R., Ott, T., and Schodel, R. 2002, MNRAS, 331,917 Gebhardt, K., Richstone, D., Tremaine, S., Lauer, T.R., Bender, R. et al. 2002,astro-ph/O209483 Genzel, R. and Townes, C.H. 1987, ARA&A, 25,377 Genzel, R., Hollenbach, D., and Townes, C.H. 1994, Rep. Prog. Phys., 57,417 Genzel, R., Thatte, N., Krabbe, A,, Kroker, H., and Tacconi-Garman, L.E. 1996, ApJ, 472, 153 Genzel, R., Pichon, C., Eckart, A., Gerhard, O.E., and Ott, T. 2000, MNRAS, 317, 348 Genzel, R., Hofmann, R., Lehnert, M., Ott, T., Schijdel, R., Eckatt, A,, Alexander, T. et al. 2003, in press Ghez, A., Klein, B.L., Moms, M., and Beckhn, E.E. 1998, ApJ, 509,678 Ghez, A., Morris, M., Becklin, E.E., Tanner, A., and Kremenek, T. 2000, Nature, 407,349 Hofmann, R., Brandl, B., Eckart, A,, Eisenhauer, F., Tacconi-Garman, L.E. 1995, Proc. SPIE, 2475, 192 Jeffries, S.M. and Christou, J.C. 1993, ApJ, 415,862 Maoz, E. 1998, ApJ, 494, L181 Melia, F. and Falcke, H. 2001, ARA&A, 39,309 Morris, M. and Serabyn, E. 1996, ARA&A, 34,645 Mouawad, N., Eckart, A., Pfalzner, S., Straubmeier, C., Spurzem, R., Genzel, R., Ott, T.,and Schodel, R. 2003, in preparation. Munyaneza, F. and Viollier, R.D. 2002, ApJ, 564,274 Ott, T., Genzel, R., Scbdel, R., and Eckart, A. 2003, in preparation Reid, M.J., Menten, K.M., Genzel, R., Ott, T., Schodel, R., and Eckart, A. 2003, ApJ, in press Schodel, R., Ott, T., Genzel, R., Hofmann, R., Lehnert, M., Eckart, A,, et al. 2002, Nature, 419,694 Schodel, R., Genzel, R., Ott,T., Eckart, A., Mouawad, N., and Alexander, T. 2003, m press Tsiklaun, D. and Viollier, R.D. 1998, ApJ, 500,591
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radius (parsec) Fig. 3 Mass distribution in the Galactic Center assuming an 8 kpc distance (Reid et al. 2003). The filled black circle denotes the mass derived from the orbit of S2, the red filled circle the mass derived from the orbit of S12, and the purple circle the mass derived from the orbit of S14. Filled dark green triangles denote Leonard-Merritt projected mass estimators from the present work (at 0.025 pc) and from a new NTT proper motion data set by Ott et al. (2003), separating late and early type stars, and correcting for the volume bias determined from Monte Carlo modeling of theoretical clusters and assuming a central density profile with a power-law slope of a = 1.37 (Genzel et al. 2003). An open down-pointing triangle denotes the Bahcall-Tremaine mass estimate obtained from Keck proper motions (Ghez et al. 1998). Light-blue, filled rectangles are mass estimates from a parameterized kanS-equdtiOn model, including anisotropy and distinguishing between late and early type stars (Genzel et al. 2000). Open circles are mass estimates from a parameterized Jeans-equation model of the radial velocities of late type stars, assuming isotropy (Genzel et al. 1996). Open red rectangles denote mass estimates from a non-parametric, maximum likelihood model, assuming isotropy and combining late and early type stars (Chakrabarty and Saha 2001). The different statistical estimates (in part using the same or similar data) agree within their uncertainties but the variations show the sensitivity to the input assumptions. In contrast, the orbital technique for S2/S12 and S14 is much simpler and less affected by the assumptions. Green letter "G" points denote mass estimates obtained from Doppler motions of gas (Genzel and Townes 1987). The blue continuous curve is the overall best fit model to all data. It is the sum of a 2.87 & 0.15 x l o 6 Ma point mass, plus the visible outer stellar cluster of central density 3.6 x 106Mopc-3, core radius 0.34 pc and power-law index a = 1.8. The grey long dash-short dash curve shows the same stellar cluster separately, but for a infinitely small core (i.e. a 'cusp'). The red dashed curve is the sum of the stellar cluster, plus a Plummer model of a hypothetical very compact (core radius -0.00019 pc) dark cluster of central density 2.2 x 10*7Mopcp3.
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Astron. Nachr./AN 324, No. S1.543-549 (2003) / DO1 10.1002/asna.200385080
Stellar Dynamics in the Galactic Center: 1000 Stars in 100 Nights Thomas Ott”, Reinhard Gemel’, Andreas Eckart2, and Rainer Schodel’ Max-Planck-Institut fkr extraterrestrische Physik, Giessenbachstr., 85748 Garching, Germany 1 . Physikalisches Institut, Universitat zu Koln, Ziilpicher Str. 77,50937 Koln, Germany
’
Key words galaxy: center, stars: infrared, stars: dynamics
Abstract. We present the results of a largely automatized re-analysis of all near-infrared imaging data obtained with the MPE speckle camera SHARP I at the NTT. We were able to increase the number of measured proper motions in the Galactic Center stellar cluster by about an order of magnitude in comparison to
previous work done on this issue. We have made astrometric positions and projected sky velocities available for about 1000 stars in the central parsec of our galaxy. We also present radial velocity measurements of 100 Stars which were obtained with the integral field spectrograph 3D at the 2.2m ESO/MPG telescope located on La Silla, Chile. Finally, using the Gemini North science verification narrow band data enables us to distinguish the stellar populations to fainter magnitudes. Using this new and extensive dataset, we present an analysis of the dynamical properties of the late-type and early-type stellar populations. While the old population is relaxed and their movements isotropic, the young stars experience noticeable anisotropy with mainly tangential proper motions.
1 Introduction Because of its proximity (distance 8 kpc, Reid 1993), the center of the Milky Way is a unique laboratory for studying the physical processes in galactic nuclei. In particular, the Galactic Center offers the unique opportunity for investigating stars and gas in the immediate vicinity of a supermassive black hole, at a level of detail that will not be accessible in any other galactic nucleus in the foreseeable future. There are several different stellar populations/components in the central parsec (for a review see Genzel 2001). The stellar mass and the near-IR light at K> 13 is dominated by red giants in the old (1-10 Gyrj component of the nuclear star cluster. A group of about two dozen luminous, blue supergiants (‘He1 emission line stars’) strongly affects the near-IR maps at the bright end (K- 9- 12j, and probably indicates recent formation of massive stars within the last 2-7 Myrs (Forrest et al. 1987, Allen, Hyland & Hillier 1990, Krabbe et al. 1991, 1995, Tamblyn et al. 1996, Blum et al. 1996, Paumard et al. 2001). A number of bright (K- 10 - 12) asymptotic giant branch (AGB) stars sample an intermediate mass, intermediate age component (2100 Myr, Lebofsky & Rieke 1987, Krabbe et al. 1995, Blum et al. 1996). Finally there is a group of dust embedded stars with near-featureless near-IR spectra (Becklin et al. 1978, Krabbe et al. 1995, Genzel et al. 1996), many of which are associated with the gaseous mini-spiral. Their nature is uncertain. The mean stellar velocities (or velocity dispersions) follow a KepIer law ((v ’) 0; R -l) from 0.1” to 2 20”, and provide compelling evidence for the presence of a central compact mass (Genzel et al. 1996, 97, 2000; Eckart and Genzel 1996,97; Ghez et al. 1998). Overall the stellar velocities are consistent with an isotropic velocity field but the He1 emission line stars appear to be preferentially on tangential orbits (Genzel et al. 2000). N
* Corresponding author: e-mail: ott@mpe,mpg.de,Phone: 4 9 89 3oooO 3276, Fax: 4 9 89 3oooO 3390 @ 2001 WILEY-VCH Verlag GmbH & Co. KGaA, Weinheirn
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2 Observations and data reduction The observations presented in this contribution were carried out using a variety of telescopes and instrumentation: The proper motion study used speckle data obtained with the MPE speckle camera SHARP I at the New Technology Telescope (NTT), which is located at La Silla, Chile and operated by the European Southern Observatory (ESO). Imaging spectroscopy using the MPE 3D spectrograph was carried out at the ESOMPG 2.2m telescope on La Silla. These data provide stellar classification and radial velocities for the brightest 100 stars in the GC. Initial astrometry and stellar identifications were obtained using NAOSlCONICA at the VLT. These data were recorded during science verification of the instrument. The public Gemini SV dataset provided stellar classification down to lower magnitudes than has been possible using the 3D data. The data reduction was carried out using the MPE dpuser software package. All data were sky subtracted, flat-fielded and deadhot pixels were corrected for by interpolation of usable neighboring pixels. In the case of the SHARP I speckle data, the final diffraction limited image was created using the Simple Shiftand-Add algorithm (see Eckart & Genzel, 1997). The Gemini and NACO data needed no further treatment. In order to calibrate the wavelength scale of the 3D data, at the beginning or the end of each observing night exposures of spectral lamps (argon lamps in our case) were done. These emit a known line-spectrum and were used to measure the dispersion of the spectrograph. The wavelength scale of the instrument was then stretched such as to get a linear relationship between detector element and wavelength. Spectral calibrator stars with a known spectrum were observed at similar airmass as the galactic center. These standard stars were divided by a spectrum of the same stellar type (Kleinmann & Hall 1986) in order to remove stellar features resulting in an atmospheric transmission spectrum. The source data were then divided by this spectrum.
3 Astrometry The task of astrometry is to determine the positions of stars for each epoch of observations. Our approach was to create a “master-list” of stars using one high-quality observation, which serves to re-identify the stars in all other observations. The measured positions then have to be transformed to a common coordinate system so they can be compared and allow the determination of proper motions. 3.1 Master-list For the creation of the master-list it is desirable to have one single observation which includes the full field-of-view of the SHARP observations, in order to eliminate systematic effects due to image distortions. 20” x 20”. Since we observed the galactic center as a mosaic, the complete field covers a field of Fortunately, in May 2002 the first observations with NAOS/CONICA at the VLT were available which cover a field of 40” x 40”, which is large enough to cover all SHARP pointings. We used the NAOSICONICA data to create a first list of stars using package “SExtractor” (Bertin & Amouts, 1996). Although SExtractor was able to identify a lot of stars, omissions in the list were common in regions of high stellar density. It was necessary to create an image in which each identified star was marked and then to manually edit the list to add sources and to remove some spurious misidentifications. This initial list of stars consists of stellar positions in pixel coordinates of the CONICA frame. It is desirable, though, to transform this to physical astrometric coordinates (relative to the radio N
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source SgrA*) in order to later derive proper motions in units of "/year or (knowing the distance to the galactic center) in k d s .
Fig. 1 Identification of masex sources in the near infrared imaging
This can be done using radio interferometry, since SgrA* is known to be a radio point source. These measurements can be done to an accuracy of a few milli-arcseconds by comparing the positions with extragalactic sources (as quasars), which is far more accurate than the positioning achievable with infrared direct imaging. In the CONICA frame, we are lucky to observe several maser sources which are both observable in the NIR and the radio (Menten et al., 1997,Reid et al., 2002). Here we used six sources (see Table 1) to calculate the plate constants of our infrared imaging. Table 1 Positions and proper motions of the maser sources relative to SgrA*
IRS 9
5.6650 f 0.0019
-6.3433 f 0.0030
IRS 7
0.0326 f 0.0030
5.5353 f 0.011
IRS 2
-3.2574 f0.0012
-6.8980 f0.0015
* 1.22 -0.91 * 0.23
SiO-B
10.4697 f0.0026
-5.8024
-0.08 f 1.30
IRS l0EE
7.6854 f 0.0010
4.2067 f0.0010
IRS 15NE
1.2197 f0.0011
11.2948
* 0.0053
* 0.0019
3.61 f 0.53
-1.13
1.72 f 0.88 -2.90
* 2.90
-2.73 f 0.28 -3.61
* 2.27
0.36 i 0.23
-2.13 3z 0.24
-1.68 f0.24
-6.04 i0.35
In order to transform the initial list of stars to this astrometric reference frame, we have calculated the centroids of the maser sources in the CONICA frame. When using six radio positions, we are able to determine image distortions up to second order. The transformation matrix was then applied to all stars in the initial list, resulting in an astrometric (relative to SgrA*) master-list of stars at the epoch of the CONICA observation.
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3.2 Identification of stars In order to determine the positions of a certain star in a given image, it is first necessary to identify this star in the image. This can be done by transforming the star’s astrometric position (2,y) (given by the masterlist) to the coordinate system (z’, y’) in the image frame. The transformation equations for the linear case are: 2’ = a2
y’ = dx
+ by + c
+ ey + f
These take into account translation, rotation, and image scale in both axes. Since it is necessary to determine six unknowns, one needs at least the positions of three stars. The greater the distance between these stars, the more reliable the solution will be. Since the SHARP I imaging was done using four mosaic settings, we needed four different triples of stars. These have been identified manually in each frame. Using the stars to Equation (1) it is then possible to transform the master-list to the current image, identifying all the stars. Of course, these positions are only accurate to within certain limits. In order to determine their positions more accurately, we have therefore searched for the maximum pixel intensity within a search radius of a few pixels. This works very well for most of the stars, only the central cluster within the central arcsecond with stars moving in excess of 500 km/s needed special treatment. Here we identified the stars manually by comparing the images of several epochs and identifying by similarity. Then we calculated the centroid ( X ,Y )of the stars in the image I ( s ,y) as follows:
As a result we then get a list in the same ordering as the master-list with the positions of the stars in the current image. Since each exposure has a different limiting magnitude G at which a reliable determination of the centroid is possible, it is in addition necessary to determine the magnitude of the faintest usable star, else the method mentioned above centers on the brightest noise peak, not resulting in a meaningful measurement of the stellar position. A stable criterion for G turned out that the brightest pixel of a star must be at least G > medzan(1mage) 0 . meddev(1mage). A reliable value for 0 was 150.
+
3.3 Transformation to astrometric coordinates When applying the steps described above for all images used in this analysis, we end up with a List of stellar positions in the detector coordinate system for each image. In order to transform them back into astrometric coordinates, a reverse transformation according to equation 1 is necessary. To do this reverse transformation more accurately and take into account image distortions of higher order, we also allow quadratic and mixed terms in the transformation matrix, which then looks like: 2’ = az
TI’= gz
+ b y + cxy
+ dz2 + ey2 + f
+ hy + i z y +j22 + ky2 + 2
(3)
This set of equations now has 12 unknowns, making it necessary to use at least six stars to solve for the unknowns. Using the six maser sources mentioned above is not possible for our data, since the SHARP I frames only have a maximum of two stars common to a single frame. Therefore we have to reference the images to the master-list given by the CONICA-observations. Our ultimate goal, though, is to derive proper motions for all the stars in our master-list. If we were to use only six stars, we would have to know
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either their proper motions beforehand, or assume that their proper motions are close to zero. Since neither of these are known, it was necessary to use a different approach: When using a large (between 50 and 200) of stars to solve for the unknowns in Eq. (3), the set of equations is, of course, over-determined. It can be solved using a least-squares approach. The only assumption about the motions of the stars then is that their average motions are close to zero, an assumption which is justified by the isotropy of the motions of most of the stars (Genzel et al., 2000). This explains also the minimum number of 50 stars necessary, since when using less stars their average motions do not cancel out using a random selection of stars. The maximum number of stars of 200 is given by the quality of our data. If using more stars, we start using measurements of centroids of faint stars with low signal-to-noise, which introduces noise in the solution of Equation (3). The results of this approach can be improved by the following method: For each star, one can estimate a transformation error by comparing its transformed position with the position given by the master-list. Stars with large errors are caused by either spurious noisy measurements, misidentifications, or extremely fast moving stars. Equation (3) can then be solved for again neglecting those stars. From the now known astrometric positions of the stars at different epochs (ranging from March 1992 to June 2002) it is straightforward to derive proper motions. For all but two stars (S1 and S2 in the central cluster, see Schodel et al., 2002) a straight line fit was sufficient to describe their motions. 3.4
Determination of radial velocities
The stellar spectra which could be extracted from the imaging spectroscopy data constitute of two different types: Stars with line emission and stars exhibiting CO bandhead absorption. In order to derive radial velocities from these stars, we used two different approaches: 0
When using stars exhibiting line-emission, one (or several in case of P-Cygni type absorption) gaussians were fit to the emission lines giving a direct measure of the redshift of this star
Stars showing CO absorption were cross-correlated with a template spectrum and the redshifts were determined from the maximum in this cross-correlation In the former case, each emission line results in an independent measure of the star’s radial velocity. These were then averaged, taking their standard deviation as an estimate for the error. In the latter case, we used the star HD 78647 as a template, taking into account its radial velocity of 9 k d s . The overall accuracy for which radial velocities could be determined ranges between 100 k d s for faint stars and better than 30 k d s for bright stars. From these measured radial velocities, we subtracted a value of 38 k d s to account for the Earth’s movement in its orbit around the sun and the movement of the solar system with respect to the GC (local frame of rest).
4 The dynamical properties of the central cluster We are considering here the normalized angular momentum along the line of sight, JZ/JZ(max),which we define as JZ/JZ(max) = (33jy
-
(4)
YV,)/PV,,
where v,, vy and up are the R.A.-, Dec.- and total proper motion velocities of a star at (x,y) on the sky -1, 0 and +1, depending on whether the stellar orbit and at projected radius p. JZ/JZ(max)is projected on the sky is mainly counter-clockwise tangential, radial or clockwise tangential with respect to the projected radius vector from the star to SgrA*. Figure 4 shows the projected radial distribution of JZ/Jz(max)for late and early type stars, as identified from the narrow-band CO-index (Figure 2). The N
N
N
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9
10
11
12
13
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K Fig. 2 CO-index (m(C0) = m(2.29)- m(2.26) ) as a function of K-magnitude for those stars in the proper motion sample that also have Gemini science demonstration data, narrow band maps. Stars marked with filled rectangles denote early type stars, and stars marked with filled circles denote late type stars confirmed by the 3D spectroscopic data. For K 5 15 stars with m ( C 0 )2 0.04 are identified as late type stars, while stars with m ( C 0 ) < 0.04 are identified as early type stars.
early type stars show a preponderance of tangential orbits (Figure 4 right panel). Within the central 3" 51(f10) % and 18(&6) % of the early type stars are on clockwise, and counter-clockwise, tangential orbits, in marked contrast to the random pattern of the late type stars (Figure 4 left panel). Early type stars with clockwise, tangential orbits dominate within a few arcseconds of SgrA*, and bunch up in the IRS16 complex WSE of SgrA".
. F
.' * I
laletypestarsK455
I-
g o >"
0
2
4
6
8
10
PmA.(-)
Fig. 3 Distribution of z-velocities of spectroscopic late type stars (squares) and early type stars (filled circles) as a function of Dee.-offset from SgrA*.
Fig. 4 Normalized angular momentum along the line of sight (J,/J,(max)) as a function of projected separation from SgrA' for the K< 15.5 late type stars ( m ( C 0 ) 2 0.04, left panel), and for the K514.7 (green) and K'/'cz 70. A few of these homothetic models have been chosen as initial guesses for other adjustments, with released constraints. It is first interesting to check the coplanar hypothesis, by keeping uniform only the two parameters that define the orbital plane, and the uniform
' Two orbits are said to be "homothetic" when they are identical except for their scale, i.e. when they share the same orbital parameters, except the periapse.
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eccentricity hypothesis. The agreement is much better when making either the eccentricity or the orbital plane free, but it is still better making both free. Even with the most general situation, the parameters are still not constrained enough to decide whether the orbits are all bound or not, to extrapolate the model outside the field of view, nor even to derive reliably the direction of proper motion. However the models share a few characteristics that we judge to be robust because of their repeatability: 1. the orbital planes are close to that of the CND; 2. the orbits are not quite coplanar; the two angles that define the orbital plane vary over a N 10”range; 3. the eccentricity varies from one orbit to another, beeing close to parabolic or above for the innermost orbits, and closer to circular (below N 0.5) for the outermost. The variations of the orbital parameters induce a particular shape for the Northem Arm (Fig. 5): for all the non-coplanar models, the Northern Arm looks like a warped surface, and this warp induces a crowding of orbits that closely follows the bright rim of the structure. That suggests that the Northern Arm is either a warped disk, or the ionized surface of a neutral cloud. The bright rim itself is not due only to the stronger UV field and a real local enhancement of the density, but also to an enhancement of the column-density due to the warp. An interesting point is that, in some models, no orbit follows the bright rim, which emphasizes that it is really important to consider the dynamics independently from the morphology of the Northern Arm. Another characteristic present in all the models is that the period of the orbits ranges from a few lo4 years to a few lo5 years, which implies that the Northern h would have a completely different shape in a few lo4 years, and cannot be much older than that timescale. Since the agreement in radial velocity is now rather good, it makes sense to look at the deviations from global motion by looking at the extended features on the residual velocity map (Fig. 5): A) the flow shows a rather significant deviation in the region just southwest of the embedded star, IRS 1W; this perturbation could be due to the interaction with this star’s wind; B) the region of this model closest to the Minicavity is perturbed; C) another deviation is seen at the precise location where the bright rim bends a lot, just east of IRS 7E2; D) finally, an elongated feature is seen on the fainter rim coming from IRS 1W towards the northeast.
4 Discussion The presence of at least three isolated gas patches (the Western Bridge, the Northern Arm Chunk and the Bar Overlay, but also possibly the Eastern Bridge, which may or may not extend ouside the field) in addition to the standard large flows has been demonstrated. In addition to that, the Microcavity at the elbow between the Eastern Arm and the Tip is a new example of an interaction between an ISM feature and a stellar wind, similar to the Minicavity, as is the deviation from Keplerian motion detected close to IRS 1W. In this context, it makes sense to ask what is the influence of the large number of mass losing stars present in the central parsec? These massive, hot stars of the central cluster, named “helium stars” from their strong 2.06 pm He I emission line, presumably LBV-type and W stars, being particularly concentrated in two clusters, IRS 16 (Krabbe et al. 1991) and IRS 13E (Maillard et al. 2003), must be a major source of helium in their environment. Therefore the following question arises: what happens to this helium enriched material? Could it form or enrich the gas patches that we see? Comparing the helium and hydrogen distribution in the central parsecs, and their abundance in the different structures and in the CND is guaranteed to help us better understand the origin of the different ISM structures. The geometry of the Northern Arm has been studied from its velocity map, leading to the conclusion that it may not be a planar structure, but rather a three-dimensional structure. Fig. 5a is quite compatible with the Northern Arm indeed being the ionized surface of a neutral cloud, which could come from the CND. in accordance with the standard formation scenario.
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10 da
0 -10 cos(6) ( Arcsec )
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10
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Fig. 5 On this Paa map (Scoville et al. 2002), one of the Keplerian models is overplotted (left panel). This one is quite consistent with the Northern Arm and the Western Arc being related structures. On the right panel, the most significant deviations from Keplerian motion discussed in the text are labeled A to D, and indicated as filled contour.
References Krabbe, A,, Genzel, R., Drapatz, S., & Rotaciuc, V. 1991, ApJ, 382, L19 Lacy, J.H., Achtemann, J.M., & Serabyn, E. 1991, ApJ, 380, L71 Lo, KY., & Clausen, M.J. 1983, Nature, 306,647 Maillard, J.P. 1995, In: Tridimensional Optical Spectroscopic Methods in Astrophysics, L4U Col. 149, G . Comte & M. Marcelin (eds), ASP Conf. Series, 71,316 Maillard, J.P. 2000, in Waging the Universe in Three Dimensions, E. van Breughel & J. Bland-Hawthorn (eds), ASP Conf. Serie 195, 185 Maillard, J.P., Paumard, T., Stolovy, S.R., & Rigaut, F. 2003, these proceedings Moms, M. & Maillard, J.P., 2000, in Imaging the Universe in Three Dimensions, E. van Breughel & J. BlandHawthorn (eds), ASP Conf. Sene 195, 196 Paumard, T., Maillard, J.P., Moms, M., & Rigaut, F. 2001, A&A, 366,466 (Paper I) Paumard, T., Maillard, J.P., Stolovy, S.R., & Rigaut, F. 2003, thesepmceedings Roberts, D.A. & Goss, W. M. 1993, ApJS, 86, 133 Scoville, N.Z., Stolovy, S.R., Rieke, M., Christopher,M., & Yusef-Zadeh, F., 2002, submitted to ApJ Vollmer, B., & Duschl, W.S. 2000, New Astronomy, 4,581 Yusef-Zadeh, F., Roberts, D.A., & Biretta, J. 1998, ApJ, 499, L159 Yusef-Zadeh, F., Stolovy, S. R., Burton, M. Wardle, M., &Ashley, M. C. B. 2001, Apl, 560,749 Zhao, J.H. & Goss, W.M. 1988, ApJ, 499, L163
Astron. Nachr./AN 324, No. SI, 613-619 (2003)/ DO1 10.1002/asna.200385095
Gas physics and dynamics in the central 50 pc of the Galaxy B. VoUmer*', W.J. Duschl***.',and R.Z ~ l k a * * ' ~
' Max-Planck-Institutfur Radioastronomie, Auf dem Hugel 69,53121 Bonn, Germany
* Institut fur Theoretische Astrophysik der Universitat Heidelberg, TiergartenstraBe 15, 69 121 Heidelberg, Germany IRAM, 300, rue de la piscine, 38406 Saint Martin dHeres, France
Key words Gas physics, gas dynamics, theory, numerical modelling Abstract. We present models for the gas physics and dynamics of the inner 50 pc of the Galaxy. In a first step the gas properties of an isolated clumpy circumnuclear disk were analytically investigated. We took the external UV radiation field, the gravitational potenQal, and the observed gas temperature into account. The model includes a demiption of the properties of individual gas clumps on small scales, and a treatment of the circumnuclear disk as a quasi-continuous accretion disk on large scales. In a second step the dynamics of an isolated circumnuclear disk were investigated with the help of a collisional N-body code. The environment of the disk is taken into account in a third step, where we calculated a pro- and a retrograde encounter of an infalling gas cloud with a pre-existing circumnuclear disk In order to constrain the dynamical model, we used the NIR absorption of the giant molecular clouds located within the inner 50 pc of the Galaxy to reconstruct their line-of-sight distribution.
1 Introduction During the discussion led by R. Narayan at this conference it became clear that the mass accretion rate onto the central black hole in the Galactic Centre at a radius of
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