FERMI/LARGE AREA TELESCOPE BRIGHT GAMMA-RAY SOURCE LIST

The Astrophysical Journal Supplement Series, 183:46–66, 2009 July  C 2009. doi:10.1088/0067-0049/183/1/46 The American Astronomical Society. All rights reserved. Printed in the U.S.A. FERMI/LARGE AREA TELESCOPE BRIGHT GAMMA-RAY SOURCE LIST A. A. Abdo1,55 , M. Ackermann2 , M. Ajello2 , W. B. Atwood3 , M. Axelsson4,5 , L. Baldini6 , J. Ballet7 , D. L. Band8,9,56 , G. Barbiellini10,11 , D. Bastieri12,13 , M. Battelino4,14 , B. M. Baughman15 , K. Bechtol2 , R. Bellazzini6 , B. Berenji2 , G. F. Bignami16 , R. D. Blandford2 , E. D. Bloom2 , E. Bonamente17,18 , A. W. Borgland2 , A. Bouvier2 , J. Bregeon6 , A. Brez6 , M. Brigida19,20 , P. Bruel21 , T. H. Burnett22 , G. A. Caliandro19,20 , R. A. Cameron2 , P. A. Caraveo23 , J. M. Casandjian7 , E. Cavazzuti24 , C. Cecchi17,18 , E. Charles2 , A. Chekhtman1,25 , C. C. Cheung9 , J. Chiang2 , S. Ciprini17,18 , R. Claus2 , J. Cohen-Tanugi26 , L. R. Cominsky27 , J. Conrad4,14,28,57 , R. Corbet9,29 , L. Costamante2 , S. Cutini24 , D. S. Davis9,29 , C. D. Dermer1 , A. de Angelis30 , A. de Luca16 , F. de Palma19,20 , S. W. Digel2 , M. Dormody3 , E. do Couto e Silva2 , P. S. Drell2 , R. Dubois2 , D. Dumora31,32 , C. Farnier26 , C. Favuzzi19,20 , S. J. Fegan21 , E. C. Ferrara9 , W. B. Focke2 , M. Frailis30 , Y. Fukazawa33 , S. Funk2 , P. Fusco19,20 , F. Gargano20 , D. Gasparrini24 , N. Gehrels9,34 , S. Germani17,18 , B. Giebels21 , N. Giglietto19,20 , P. Giommi24 , F. Giordano19,20 , T. Glanzman2 , G. Godfrey2 , I. A. Grenier7 , M.-H. Grondin31,32 , J. E. Grove1 , L. Guillemot31,32 , S. Guiriec35 , Y. Hanabata33 , A. K. Harding9 , R. C. Hartman9 , M. Hayashida2 , E. Hays9 , S. E. Healey2 , D. Horan21 , R. E. Hughes15 , G. Jóhannesson2 , A. S. Johnson2 , R. P. Johnson3 , T. J. Johnson9,34 , W. N. Johnson1 , T. Kamae2 , H. Katagiri33 , J. Kataoka36 , N. Kawai37,38 , M. Kerr22 , J. Knödlseder39 , D. Kocevski2 , M. L. Kocian2 , N. Komin7,26 , F. Kuehn15 , M. Kuss6 , J. Lande2 , L. Latronico6 , S.-H. Lee2 , M. Lemoine-Goumard31,32 , F. Longo10,11 , F. Loparco19,20 , B. Lott31,32 , M. N. Lovellette1 , P. Lubrano17,18 , G. M. Madejski2 , A. Makeev1,25 , M. Marelli23 , M. N. Mazziotta20 , W. McConville9,34 , J. E. McEnery9 , S. McGlynn4,14 , C. Meurer4,28 , P. F. Michelson2 , W. Mitthumsiri2 , T. Mizuno33 , A. A. Moiseev8,34 , C. Monte19,20 , M. E. Monzani2 , E. Moretti10,11 , A. Morselli40 , I. V. Moskalenko2 , S. Murgia2 , T. Nakamori38 , P. L. Nolan2 , J. P. Norris41 , E. Nuss26 , M. Ohno42 , T. Ohsugi33 , N. Omodei6 , E. Orlando43 , J. F. Ormes41 , M. Ozaki42 , D. Paneque2 , J. H. Panetta2 , D. Parent31,32 , V. Pelassa26 , M. Pepe17,18 , M. Pesce-Rollins6 , F. Piron26 , T. A. Porter3 , L. Poupard7 , S. Rainò19,20 , R. Rando12,13 , P. S. Ray1 , M. Razzano6 , N. Rea44,45 , A. Reimer2 , O. Reimer2 , T. Reposeur31,32 , S. Ritz9 , L. S. Rochester2 , A. Y. Rodriguez45 , R. W. Romani2 , M. Roth22 , F. Ryde4,14 , H. F.-W. Sadrozinski3 , D. Sanchez21 , A. Sander15 , P. M. Saz Parkinson3 , J. D. Scargle46 , T. L. Schalk3 , A. Sellerholm4,28 , C. Sgrò6 , M. S. Shaw2 , C. Shrader8 , A. Sierpowska-Bartosik45 , E. J. Siskind47 , D. A. Smith31,32 , P. D. Smith15 , G. Spandre6 , P. Spinelli19,20 , J.-L. Starck7 , T. E. Stephens46,48 , M. S. Strickman1 , A. W. Strong43 , D. J. Suson49 , H. Tajima2 , H. Takahashi33 , T. Takahashi42 , T. Tanaka2 , J. B. Thayer2 , J. G. Thayer2 , D. J. Thompson9 , L. Tibaldo12,13 , O. Tibolla50 , D. F. Torres45,51 , G. Tosti17,18 , A. Tramacere2,52 , Y. Uchiyama2 , T. L. Usher2 , A. Van Etten2 , N. Vilchez39 , V. Vitale40,53 , A. P. Waite2 , E. Wallace22 , P. Wang2 , K. Watters2 , B. L. Winer15 , K. S. Wood1 , T. Ylinen4,14,54 , M. Ziegler3 (The Fermi/LAT Collaboration) 2 1 Space Science Division, Naval Research Laboratory, Washington, DC 20375, USA W. W. Hansen Experimental Physics Laboratory, Kavli Institute for Particle Astrophysics and Cosmology, Department of Physics and SLAC National Accelerator Laboratory, Stanford University, Stanford, CA 94305, USA; [email protected] 3 Santa Cruz Institute for Particle Physics, Department of Physics and Department of Astronomy and Astrophysics, University of California at Santa Cruz, CA 95064, USA 4 The Oskar Klein Centre for Cosmo Particle Physics, AlbaNova, SE-106 91 Stockholm, Sweden 5 Department of Astronomy, Stockholm University, SE-106 91 Stockholm, Sweden 6 Istituto Nazionale di Fisica Nucleare, Sezione di Pisa, I-56127 Pisa, Italy 7 Laboratoire AIM, CEA-IRFU/CNRS/Université Paris Diderot, Service d’Astrophysique, CEA Saclay, F-91191 Gif sur Yvette, France; [email protected], [email protected] 8 Center for Research and Exploration in Space Science and Technology (CRESST), NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA 9 NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA; [email protected] 10 Istituto Nazionale di Fisica Nucleare, Sezione di Trieste, I-34127 Trieste, Italy 11 Dipartimento di Fisica, Università di Trieste, I-34127 Trieste, Italy 12 Istituto Nazionale di Fisica Nucleare, Sezione di Padova, I-35131 Padova, Italy 13 Dipartimento di Fisica “G. Galilei,” Università di Padova, I-35131 Padova, Italy 14 Department of Physics, Royal Institute of Technology (KTH), AlbaNova, SE-106 91 Stockholm, Sweden 15 Department of Physics, Center for Cosmology and Astro-Particle Physics, The Ohio State University, Columbus, OH 43210, USA 16 Istituto Universitario di Studi Superiori (IUSS), I-27100 Pavia, Italy 17 Istituto Nazionale di Fisica Nucleare, Sezione di Perugia, I-06123 Perugia, Italy 18 Dipartimento di Fisica, Università degli Studi di Perugia, I-06123 Perugia, Italy 19 Dipartimento di Fisica “M. Merlin” dell’Università e del Politecnico di Bari, I-70126 Bari, Italy 20 Istituto Nazionale di Fisica Nucleare, Sezione di Bari, I-70126 Bari, Italy 21 Laboratoire Leprince-Ringuet, École polytechnique, CNRS/IN2P3, F-91128 Palaiseau, France 22 Department of Physics, University of Washington, Seattle, WA 98195-1560, USA 23 INAF-Istituto di Astrofisica Spaziale e Fisica Cosmica, I-20133 Milano, Italy 24 Agenzia Spaziale Italiana (ASI) Science Data Center, I-00044 Frascati (Roma), Italy 25 George Mason University, Fairfax, VA 22030, USA 26 Laboratoire de Physique Théorique et Astroparticules, Université Montpellier 2, CNRS/IN2P3, F-34095 Montpellier, France 27 Department of Physics and Astronomy, Sonoma State University, Rohnert Park, CA 94928-3609, USA 28 Department of Physics, Stockholm University, AlbaNova, SE-106 91 Stockholm, Sweden 29 University of Maryland, Baltimore County, Baltimore, MD 21250, USA 30 Dipartimento di Fisica, Università di Udine and Istituto Nazionale di Fisica Nucleare, Sezione di Trieste, Gruppo Collegato di Udine, I-33100 Udine, Italy 31 CNRS/IN2P3, Centre d’Études Nucléaires Bordeaux Gradignan, UMR 5797, F-33175 Gradignan, France 46 No. 1, 2009 FERMI/LAT BRIGHT SOURCE LIST 47 32 Université de Bordeaux, Centre d’Études Nucléaires Bordeaux Gradignan, UMR 5797, F-33175 Gradignan, France 33 Department of Physical Sciences, Hiroshima University, Higashi-Hiroshima, Hiroshima 739-8526, Japan 34 University of Maryland, College Park, MD 20742, USA 35 University of Alabama in Huntsville, AL 35899, USA 36 Waseda University, 1-104 Totsukamachi, Shinjuku-ku, Tokyo 169-8050, Japan 37 Cosmic Radiation Laboratory, Institute of Physical and Chemical Research (RIKEN), Wako, Saitama 351-0198, Japan 38 Department of Physics, Tokyo Institute of Technology, Meguro City, Tokyo 152-8551, Japan 39 Centre d’Étude Spatiale des Rayonnements, CNRS/UPS, BP 44346, F-30128 Toulouse Cedex 4, France 40 Istituto Nazionale di Fisica Nucleare, Sezione di Roma “Tor Vergata,” I-00133 Roma, Italy 41 Department of Physics and Astronomy, University of Denver, Denver, CO 80208, USA 42 Institute of Space and Astronautical Science, JAXA, 3-1-1 Yoshinodai, Sagamihara, Kanagawa 229-8510, Japan 43 Max-Planck Institut für extraterrestrische Physik, D-85748 Garching, Germany 44 Sterrenkundig Institut “Anton Pannekoek,” NL-1098 SJ Amsterdam, Netherlands 45 Institut de Ciencies de l’Espai (IEEC-CSIC), Campus UAB, E-08193 Barcelona, Spain 46 Space Sciences Division, NASA Ames Research Center, Moffett Field, CA 94035-1000, USA 47 NYCB Real-Time Computing Inc., Lattingtown, NY 11560-1025, USA 48 Universities Space Research Association (USRA), Columbia, MD 21044, USA 49 Department of Chemistry and Physics, Purdue University Calumet, Hammond, IN 46323-2094, USA 50 Max-Planck-Institut für Kernphysik, D-69029 Heidelberg, Germany 51 Institució Catalana de Recerca i Estudis Avançats (ICREA), E-08010 Barcelona, Spain 52 Consorzio Interuniversitario per la Fisica Spaziale (CIFS), I-10133 Torino, Italy 53 Dipartimento di Fisica, Università di Roma “Tor Vergata,” I-00133 Roma, Italy 54 School of Pure and Applied Natural Sciences, University of Kalmar, SE-391 82 Kalmar, Sweden Received 2009 February 8; accepted 2009 May 19; published 2009 June 16 ABSTRACT Following its launch in 2008 June, the Fermi Gamma-ray Space Telescope (Fermi) began a sky survey in August. The Large Area Telescope (LAT) on Fermi in three months produced a deeper and better resolved map of the γ -ray sky than any previous space mission. We present here initial results for energies above 100 MeV for the 205 most significant (statistical significance greater than ∼10σ ) γ -ray sources in these data. These are the best characterized and best localized point-like (i.e., spatially unresolved) γ -ray sources in the early mission data. Key words: galaxies: active – gamma rays: observations – pulsars: general – surveys 1. INTRODUCTION Collections of information about what can be seen in the sky range from simple lists to complex catalogs. For highenergy γ -rays (photon energies above 100 MeV), the first effort of this type was a COS-B source list (Hermsen et al. 1977), followed by the second COS-B catalog (Swanenburg et al. 1981). The Energetic Gamma Ray Experiment Telescope (EGRET) on the Compton Gamma Ray Observatory yielded several catalogs, culminating in the third EGRET Catalog (3EG; Hartman et al. 1999) and an alternate catalog, EGR (Casandjian & Grenier 2008), but also including a catalog of just the sources seen above 1 GeV (Lamb and Macomb, 1997). The AGILE telescope has recently released its first catalog (Pittori et al. 2009).58 The rapidly changing field of TeV γ -ray astronomy has a number of online catalogs, e.g., TeVCat,59 a frequently updated compilation of announced TeV sources from groundbased observatories. The Fermi Gamma-ray Space Telescope (Fermi) Large Area Telescope (LAT) is a successor to EGRET, with greatly improved sensitivity, angular resolution, and energy range. This paper presents a list of bright LAT sources that have statistical significances of 10σ or higher, based on the first three months of survey data. Although the first official LAT catalog is planned for release after the first year of operations (after the LAT gamma55 National Research Council Research Associate. Deceased. 57 Royal Swedish Academy of Sciences Research Fellow, funded by a grant from the K. A. Wallenberg Foundation. 58 See http://www.asdc.asi.it/agilebrightcat/. 59 http://tevcat.uchicago.edu/. 56 ray data themselves become publicly available),60 this early list of bright sources was released to enable multiwavelength studies by the broader community and to support proposal preparation for Cycle 2 of the Fermi Guest Investigator program. The reader is cautioned to avoid generalizing from this sample of sources. Some particular features are as follows: 1. The source list is not a complete summary of sources seen by the LAT. Many additional sources are detected with lower confidence levels in the LAT data than are included here (Section 3.3). 2. The source list is not flux limited and hence not uniform. Only sources above a 10σ statistical significance are included, as described below. Moreover, owing to the strong energy dependence both of the angular resolution of the LAT and of the intensities of backgrounds, the limiting flux is dependent on spectral hardness. Because γ -ray sources are seen against a background of diffuse gamma radiation, which is highly nonuniform across the sky, e.g., Hunter et al. (1997) and Strong et al. (2004), the limiting flux for a given statistical significance and spectral shape varies with position (Section 3.3). 3. The source list does not include detailed information about the energy spectra of individual sources. Because this list is a step toward the first LAT catalog, we adopt the terminology for sources that will be used in that catalog, with a 0 prefix. The source designation is 0FGL JHHMM.m+DDMM where the 0 refers to the preliminary nature of this list and FGL represents Fermi Gamma-ray LAT (Section 5). 60 See http://fermi.gsfc.nasa.gov/ssc/proposals/. 48 ABDO ET AL. 2. GAMMA-RAY DETECTION WITH THE LARGE AREA TELESCOPE The LAT is a pair-production telescope (Atwood et al. 2009). The tracking section has 36 layers of silicon microstrip detectors to record the tracks of charged particles, interleaved with 16 layers of tungsten foil (12 thin layers, 0.03 radiation length, at the top or front of the instrument, followed by 4 thick layers, 0.18 radiation length, in the back section) to promote γ -ray pair conversion. Below the tracker lies an array of CsI crystals to determine the γ -ray energy. The tracker is surrounded by segmented charged-particle anticoincidence detectors (plastic scintillators with photomultiplier tubes) to reject cosmic-ray backgrounds. The LAT’s improved sensitivity compared to EGRET stems from a large peak effective area (∼8000 cm2 , or ∼6 times greater than EGRET’s), large field of view (∼2.4 sr, or nearly five times greater than EGRET’s), good background rejection, superior angular resolution (68% containment angle ∼ 0.◦ 6 at 1 GeV for the front section and about a factor of 2 larger for the back section versus ∼1.◦ 7 at 1 GeV for EGRET; Thompson et al. 1993), and improved observing efficiency (keeping the sky in the field of view with scanning observations versus inertial pointing for EGRET). Pre-launch predictions of the instrument performance are described in Atwood et al. (2009). Verification of the on-orbit response is in progress (Abdo et al. 2009q) but the indications are that it is close to expectations. The data analyzed for this source list were obtained during 2008 August 4–2008 October 30 (LAT runs 239503624– 247081608, where the numbers refer to the Mission Elapsed Time, or MET, in seconds since 00:00 UTC on 2001 January 1). During this time Fermi was operated in the sky scanning survey mode (viewing direction rocking 35◦ north and south of the zenith on alternate orbits), except for a few hours of special calibration observations during which the rocking angle was much larger than nominal for the survey mode or the configuration of the LAT was different from normal for science operations. Time intervals when the rocking angle was larger than 47◦ have been excluded from the analysis because the bright limb of the Earth enters the field of view (see below). In addition, two short time intervals associated with gamma-ray bursts (GRBs) that were detected in the LAT have been excluded. These intervals correspond to GRB 080916C (MET 243216749–243217979; (Abdo et al. 2009a)) and GRB 081024B (MET 246576157– 246576187). The total live time included is 7.53 Ms, corresponding to 82% efficiency after accounting for readout dead time and for observing time lost to passages through the South Atlantic Anomaly (∼13%). The standard onboard filtering, event reconstruction, and classification were applied to the data (Atwood et al. 2009), and for this analysis the “Diffuse” event class61 is used. This is the class with the least residual contamination from chargedparticle backgrounds. The tradeoff for using this event class is primarily reduced effective area, especially below 500 MeV. Test analyses were made with the looser “Source” class cuts and these were found to be less sensitive overall than the Diffuse class for source detection and characterization. The alignment of the Fermi observatory viewing direction with the z-axis of the LAT was found to be stable during survey mode observation (Abdo et al. 2009q). The instrument response functions—effective area, energy redistribution, and 61 See http://fermi.gsfc.nasa.gov/ssc/data/analysis/documentation/Cicerone/ Cicerone_Data/LAT_DP.html. Vol. 183 point-spread function (PSF)—used in the likelihood analyses described below were derived from GEANT4-based Monte Carlo simulations of the LAT using the event selections corresponding to the Diffuse event class. The Monte Carlo simulations themselves were calibrated prior to launch using accelerator tests of flight-spare “towers” of the LAT (Atwood et al. 2009). Consistency checks with observations of bright sources in flight data are in progress (Abdo et al. 2009q). Early indications are that the effective area below 100 MeV was overestimated by as much as 30% owing to pile-up effects in the detectors. The source detection and spectral fitting analyses described below use only data >200 MeV. The impact of the lower-than-predicted effective area below 200 MeV is limited. The Diffuse event class already had relatively little effective area below 200 MeV, and so the impact on sensitivity for source detection is small. Analyses of flight data suggest that the PSF is somewhat broader than the calculated Diffuse class PSF at high energies; the primary effect for the current analysis is to decrease the localization capability somewhat. For the bright source analysis a cut on zenith angle was applied to the Diffuse class events to limit the contamination from albedo γ -rays from interactions of cosmic rays with the upper atmosphere of the Earth. These interactions make the limb of the Earth (zenith angle ∼113◦ at the 565 km, nearly circular orbit of Fermi) an intensely bright γ -ray source (Thompson et al. 1981). The limb is very far off-axis in survey mode observations, but during a small fraction of the time range included in this analysis the rocking angle reached angles as great as 47◦ (see above) and so the limb was only ∼66◦ off-axis. Removing events at zenith angles greater than 105◦ affects the exposure calculation negligibly but reduces the overall background rate. After these cuts, the data set contains 2.8 × 106 γ -rays with energies >100 MeV. Figures 1 and 2 summarize the data set used for this analysis. The intensity map of Figure 1 shows the dramatic increase at low Galactic latitudes of the brightness of the γ -ray sky. Figure 2 shows the corresponding exposure map for the representative energy 1 GeV. The average exposure is ∼1 Ms and nonuniformities are relatively small (about 30% difference between minimum and maximum), with the deficit around the south celestial pole due to loss of exposure during passages of Fermi through the South Atlantic Anomaly (Atwood et al. 2009). 3. CONSTRUCTION OF THE BRIGHT SOURCE LIST Although Figure 1 shows some obvious bright sources, finding and measuring the properties of even the high-confidence sources involves more than visual inspection of the map. Because this analysis involves the entire sky and the broad energy range of the LAT, it is necessarily more complex than the analysis of an individual source. The source list was built on the basis of the full time interval. That is, we did not attempt to detect potentially flaring sources on shorter timescales, although we did check for variability of the sources (Section 3.5) after the list was constructed. Three steps were applied in sequence: detection, localization, significance estimate. At each step only a subset of the list at the previous step was kept. In that scheme the bright source list threshold is defined at the last step, but the completeness is controlled by the first one. After the list was defined we determined the source characteristics (flux in two energy bands, time variability) and we searched for possible counterparts. No. 1, 2009 FERMI/LAT BRIGHT SOURCE LIST 49 Figure 1. Sky map of the LAT data for the time range analyzed in this paper, Aitoff projection in Galactic coordinates. The image shows gamma-ray intensity for energies >300 MeV in units of photons m−2 s−1 sr−1 . Figure 2. Exposure of the LAT for the time range analyzed in this paper, Aitoff projection in Galactic coordinates. The units are equivalent on-axis exposure in Ms. 3.1. Detection At this time we do not have a good way to look for sources directly in the three-dimensional space of position and energy so we used standard image detection techniques on counts images integrated over energy, in which each event is simply stacked into the pixel corresponding to its best-guess incident direction. The algorithm we used (mr_filter) is based on the wavelet analysis in the Poisson regime (Starck & Pierre 1998). It looks for local deviations from the background model, leaving the background normalization free. We used the same background model defined in Section 3.3, but without any spectral correction. It returns a map of significant features (above some threshold) on which we run a peak-finding algorithm, SExtractor (Bertin & Arnouts 1996), to end up with a list of sources. We also used for comparison another wavelet algorithm (PGWave, Damiani et al. 1997; Ciprini et al. 2007) which differs in the detailed implementation and returns directly a list of sources. Pre-launch simulations have shown that the latter was somewhat more sensitive on a flat background (i.e., at high Galactic latitudes) but did not work as well in the Galactic plane. At the 10σ level, the two detection methods yield identical source lists. An important decision was which energy bands to use when applying the detection algorithms. The most important instrumental characteristic in this respect is the PSF. The 68% containment radius improves by a factor of 25, from ∼ 5◦ at 100 MeV to better than 1◦ at 1 GeV, reaching ∼ 0.◦ 2 above 10 GeV (Atwood et al. 2009). For this reason there is (at least over three months) little confusion above 1 GeV and the diffuse background is not very limiting except in the Galactic ridge. On the other hand, most of the photons (83%) are recorded below 1 GeV. The majority of the sources in the Galactic plane have overlapping PSFs and are background dominated below 1 GeV (i.e., there are more background than source events inside the PSF). The starting energy therefore represents a trade between statistics and resolution. Another important aspect is that the events converted in the top, thin layers of the tracker (Front events) have nearly a factor of 2 better PSF at a given energy than those converted in the 50 ABDO ET AL. bottom thick layers (Back events). This corresponds to Back events of energy E having the same PSF width as Front events of energy E/2. Therefore, to optimize the sensitivity of the source detection we used separate energy selections for Front and Back events. The final scheme combines three energy bands. The full detection band (1.8 × 106 events) starts at 200 MeV for Front and 400 MeV for Back events. The remainder of the 2.8 × 106 events above 100 MeV carry little position information and were not used for source detection. We use a medium band starting at 1 GeV for Front and 2 GeV for Back events (3.2 × 105 events), which provides better position estimates for hard spectrum sources. We have also used a high-energy band starting at 5 GeV for Front and 10 GeV for Back events. This band is very photon starved (3 × 104 events) but has essentially no background in a PSF-sized region and can be useful for very hard sources and to avoid confusion in the Galactic plane. We use smaller image pixels at high energy (0.◦ 1) than in the medium band (0.◦ 2) and the full energy band (0.◦ 3) to adapt to the broader PSF at low energy. The bands are not exclusive (i.e., the full band includes the highenergy photons) because the high-energy events always improve the detection. To obtain a global list of candidate sources we start with the sources detected in the high-energy band (best localization) and add the sources detected in the lower-energy bands in turn, excluding sources whose positions are consistent with detections at higher energies. Because the source detection methods are standard algorithms not specific to Fermi they work in Cartesian coordinates, not the spherical sky. We map the whole sky with 24 local World Coordinate System projections (Calabretta & Greisen 2002) in Galactic coordinates: four CAR (plate carrée) projections along the Galactic plane covering −10◦ to +10◦ , six AIT (HammerAitoff) projections on each side of the plane covering 10◦ –45◦ , and four ARC (zenithal equidistant) projections (rotated 45◦ so that the pole is in a corner) covering 45◦ –90◦ . Each map is 5◦ larger on each side than the area from which the sources are extracted, to avoid border effects. We set the threshold of the source detection step at 4σ . This resulted in 562 “seed” sources. 290 were best detected in the full band, 212 in the medium band, and 60 in the high band (among 151 total excesses above 4σ in that band). 3.2. Localization The image-based detection algorithms provide estimates of the source positions, but the positions are not optimal because the energy-dependent extent of the PSF is not fully taken into account. These methods also do not supply error estimates on the positions. The method that we use to localize the sources (pointfit) is a binned likelihood technique. It uses relatively narrow energy bins (typically four per decade) and sums log(likelihood) over the energy bins. It does not use events below 500 MeV, which carry little information on position. To optimize the technique further the analysis gathers Front and Back events according to their PSF widths rather than their energies. Each source is treated independently. This means the model is a point source (with the same position but different width in each energy bin following the PSF) on top of a background model with free scaling in each energy bin. The sources are treated in descending order such that brighter sources are included in the background model for fainter ones. The closest nearby source was 0.◦ 5, with only a small effect on the fits to the lower-energy bins. The program returns the best-fit position and the error estimate (1σ along one Vol. 183 dimension) based on the assumption that −2∆log(likelihood) behaves as a χ 2 distribution. The LAT PSF itself is very close to axisymmetric (Atwood et al. 2009). The error box is not in general circular due to fluctuations in the positions of the few high-energy photons that dominate the localization precision. Here we neglect this effect, which is small for strong sources, and provide only error circles. Of the 562 initial sources, pointfit did not converge for 50 at this step, or converged to another nearby source. The reason could be confusion, or a soft spectrum leading to too few source events above 500 MeV. We did not discard these outright, but kept their original positions. Several of them were deemed significant by the maximum likelihood algorithm (Section 3.3). We defined the positions and position uncertainties of these using a more precise but much slower tool (gtfindsrc) which accounts for all sources in the vicinity. More precisely, we included in the local model all nearby sources (even those below the bright source limit defined in Section 3.3). The spectral parameters of those within 1◦ of the current source were left free, but only the current source’s position was adjusted in a given run. The same tool was used in a number of confused regions (mostly close to the Galactic plane) in which the primary analysis did not converge well. Thirty-two sources in all were treated that way, including 13 of the bright sources presented here. In the end 532 sources survived the localization step. The angular uncertainties for localization are determined from the shape of the likelihood function as described above. This results in a one-dimensional 1σ error estimate ∆xstat . For a two-dimensional axisymmetric Gaussian distribution the 95%confidence level radius r95 is related to ∆xstat by a factor −2 log(1 − 0.95) = 2.45. However, examining the distribution of the position errors from high-confidence, identified sources, we found that we needed to increase the uncertainties by 40% in order to be sure of including 95% of the cases. For very bright sources like Vela, the observed offsets from the true position observed with the present analysis led us to add in quadrature an additional systematic uncertainty of 0.04 deg to r95 : 2 r95 = (1.4 × 2.45 × ∆xstat )2 + (0.◦ 04)2 . (1) Both the 1.4 correction factor and the 0.◦ 04 systematic uncertainty are conservative and are expected to improve. Figure 3 illustrates the resulting position uncertainties as a function of the Test Statistic (TS) values obtained in Section 3.3. The relatively large dispersion that is seen at a given TS is in part due to the different local conditions (level of diffuse γ -ray emission) but primarily to the source spectrum. Hard sources are better localized than soft ones for the same TS because the PSF is so much narrower at high energy. At our threshold of TS = 100 (10σ ) the typical 95% uncertainty radius is about 10′ and the maximum is 20′ . 3.3. Significance and Thresholding The detection and localization steps provide estimates of significance, but these are underestimates because the detection step does not explicitly use the energy information and the localization step does not use the low-energy events. To better estimate the source significances we use a three-dimensional maximum likelihood algorithm (gtlike) in the unbinned mode, i.e., each event is considered individually according to its direction, energy, and conversion location in the LAT. This is part of the standard Science Tools software package62 currently 62 http://fermi.gsfc.nasa.gov/ssc/data/analysis/documentation/Cicerone/. No. 1, 2009 FERMI/LAT BRIGHT SOURCE LIST Figure 3. Source location uncertainty radii (r95 from Equation (1)) as a function of √ Test Statistic (Section 3.3), down to a limit of TS = 25. The dotted line is a 1/ TS trend for reference. The vertical dashed line is our TS = 100 threshold. The horizontal dashed line is the absolute systematic error that we adopted. at version 9r9. The gtlike tool provides for each source the bestfit parameters and the Test Statistic TS = 2∆log(likelihood) between models with and without the source. This tool does not vary the source position, but it adjusts the source spectrum. It should be noted that gtlike does not include the energy dispersion in the TS calculation (i.e., it assumes that the measured energy is the true energy). Given the 8%–10% 1σ energy resolution of the LAT over the energy bands used in the present analyses, this approximation is justified. The underlying optimization engine is Minuit.63 The code works well with up to ∼30 free parameters, an important consideration for regions where sources are close enough together to partially overlap. Uncertainty estimates (and a full covariance matrix) are obtained from Minuit in the quadratic approximation around the best fit. For this stage we modeled the sources with simple power-law spectra. The TS associated with each source is a measure of the source significance, or equivalently the probability that such an excess can be obtained from background fluctuations alone. The probability distribution in such a situation (source over background) is not precisely known (Protassov et al. 2002). However, since we consider only positive fluctuations, and each fit involves 2 degrees of freedom (flux and spectral slope), the probability to get at least TS at a given position in the sky is close to 1/2 of the χ 2 distribution with 2 degrees of freedom (Mattox et al. 1996), so that TS = 25 corresponds to 4.6σ (one sided). Pre-launch simulations have shown that this approximation is indeed true if the background model is close to the truth. The diffuse background is of course very important since it represents around 90% of the events. We model the Galactic diffuse emission using GALPROP, described in Strong et al. (2004) and Strong (2007), which uses a realistic representation of cosmic-ray propagation in the Galaxy and the resulting γ -ray emission; it uses distributions of gas based on radioastronomical surveys, and the interstellar radiation field from an extensive modeling package. For this work, the GALPROP package has been updated to include recent H i and CO surveys, more accurate decomposition into Galactocentric rings, as well as a new calculation of the interstellar radiation field for inverse Compton emission (Porter et al. 2008). For this work the fit of 63 http://lcgapp.cern.ch/project/cls/work-packages/mathlibs/minuit/ doc/doc.html. 51 the model to the Fermi data was improved by an increase in the inverse Compton component, and a flatter cosmic-ray gradient in the outer Galaxy. The particular GALPROP run designation for our model is 54− 59varh7S. Because the fitted fluxes and spectra of the sources can be very sensitive to even slight errors in the spectral shape of the diffuse emission we allow the Galactic diffuse model to be corrected (i.e., multiplied) locally by a power law in energy with free normalization and spectral slope. The slope varies between 0 and 0.15 (making it harder) in the Galactic plane and the normalization by ± 20%. The isotropic component of the diffuse emission represents the extragalactic and residual backgrounds (instrumental + Earth albedo). It is modeled by a simple power law. Its spectral slope was fixed to E −2.25 , the best-fit value at high latitude, and its normalization was left free. The three free parameters were separately adjusted in each Region of Interest (RoI; see below). For this significance analysis, we used only events with energies above 200 MeV because the fits to the diffuse spectrum were systematically high below 200 MeV; the extrapolation of the high-energy spectrum overestimated the data, possibly because of the acceptance bias described in Section 3.6. We feared that including the low-energy points could bias the whole process. This energy cut changes little the TS estimates except for the very softest sources. The high energy limit for the analysis was set to 100 GeV. There were fewer than 1000 events above 100 GeV, and at this point we do not have a single source that is bright enough to check our calibration above that limit. We split the sky into overlapping circular RoI, each typically 15◦ in radius. The source parameters are free in the central part of each RoI (which is chosen such that all free sources are well within the RoI even at low energy). We adjust the RoI size so that not more than eight sources are free at a time. Adding three parameters for the diffuse model, the total number of free parameters in each RoI is 19 at most. We needed 128 RoIs to cover the 532 seed positions. We proceed iteratively. All RoIs are processed in parallel and a global current model is assembled after each step in which the best-fit parameters for each source are taken from the RoI whose center is closest to the source. At each step the parameters of the sources close to the borders are fixed to their values in the global model at the end of the previous step; they all start at 0 flux at the first step (the starting point for the spectral slope is 2). Sources formally outside the RoI (but which can contribute at low energy due to the broad PSF) are included in the model as well. We iterate over five steps (the fits change very little after the fourth). At each step we remove sources with low TS and refit, raising the threshold up to 25 (approximately 5σ ) at the last step. We have checked via simulations that removing the faint sources has less impact on the bright sources than does changing the diffuse model (Section 3.6). This procedure left 444 sources, among which 205 have TS > 100. We chose not to include the lower-significance sources (TS < 100) in the bright source list for the following two reasons. 1. The number of sources per TS interval normally decreases with increasing TS for any log N–log S close to Euclidean. This is not the case with our procedure (there are fewer sources at 25 < TS < 30 than at 35 < TS < 40), particularly in the Galactic plane. This is a rather sure sign that we are missing sources at low flux, and more so in the Galactic plane. Given the relatively rough nature of the 52 ABDO ET AL. Vol. 183 -6 ×10 0.4 0.35 0.3 0.25 0.2 0.15 0.1 0.05 100 MeV in photons cm−2 s−1 ) needed for a 10σ Figure 4. Flux (E > photon spectral index is 2.2. Galactic coordinates. detection for the LAT data for the three month time range considered in this paper. The assumed detection procedure (Section 3.1) this is not particularly surprising. 2. Judging by the spatial and spectral residuals the Galactic diffuse model is still in need of improvement (see Section 3.6). This uncertainty makes us wary of claiming detections of sources not too far above the diffuse level. On the other hand the sources at TS > 100 can be seen by eye on the images and we are confident they are all real. Note that all excesses formally above TS = 25 were included in the maximum likelihood adjustments (including those described in Sections 3.4 and 3.5) to avoid transferring their fluxes to the more-significant sources. Figure 4 shows the source flux needed to reach a 10σ significance level at any point in the sky for the three month time interval considered in this analysis. This is based on a calculation using the Galactic diffuse and isotropic background models, the instrument response functions of the LAT, and the pointing history during the three months, and assuming an E −2.2 spectrum, the average spectral shape of the sources. This should be viewed as an indication only because the detection threshold depends on the source spectrum. Although the nonuniform exposure affects this map somewhat, the dominant factor is the strong diffuse emission along the Galactic plane. 3.4. Flux Determination The maximum likelihood method described in Section 3.3 provides good estimates of the source significances, but not very accurate estimates of the fluxes. This is because the spectra of most sources do not follow a single power law over that broad an energy range (more than two decades). Among the two most populous classes, the active galactic nuclei (AGNs) often show a broken power-law spectrum and the pulsars an exponentially cutoff power law. In both cases, fitting a single power law over the entire range overshoots at low energy where most of the photons are, and therefore biases the fluxes high (on the other hand the effect on the significance is low due to the broad PSF and high background at low energies). An additional difficulty is that the fit over the entire range stopped at 200 MeV, whereas comparison with previous missions requires that we provide fluxes starting at 100 MeV. Extrapolating back to 100 MeV would have added another error. To provide better estimates of the source fluxes, we have decided to split the range into two and define two independent bands from 100 MeV to 1 GeV and 1 GeV to 100 GeV. The 1 GeV limit is largely arbitrary but is a round number that happens to split the data into approximately equal contributions to the sources’ significance. The list of sources remains the same in the two bands of course. Each band is treated in the same way as the full band in Section 3.3. The power-law slopes are fitted independently in each band for each source. We discard the slopes here because they are not very precise (the low band is not very broad and there are not many events in the high band) and keep only the flux estimates. Even though the fit is not good near 100 MeV as mentioned in Section 3.3 (see also Section 3.6), including the data down to 100 MeV still provides a more reliable estimate of the flux than extrapolation for all sources which do not follow exactly a power law. The estimate from the sum of the two bands is on average within 30% of the flux obtained in the previous section, with excursions up to a factor of 2. We have also compared these estimates with a more precise spectral model for the three bright pulsars (Vela, Geminga, and the Crab). The flux estimates are within 5% of each other. An additional difficulty that does not exist when considering the full data is that, because we wish to provide the fluxes in both bands for all sources, we must handle the case of sources that are not significant in one of the bands or where the flux is poorly determined due to large uncertainty in the spectrum. This situation occurs even for the high-confidence sources reported here: nine have TS < 25 for the 100 MeV to 1 GeV band, and two have TS < 10 in this band. No high-confidence source has TS < 25 in the 1–100 GeV band. This difference reflects the fact that the current study (and the LAT in general) is more sensitive at high energy. For the sources with TS < 10 or poorly measured flux values (where the nominal uncertainty is comparable to the flux itself), we replace the flux value from the likelihood analysis by a 2σ upper limit (2∆log(likelihood) = 4), indicating the upper limit by a 0 in the flux uncertainty column of the source list table. 3.5. Variability For this paper, we wanted to flag sources that are clearly variable. To that end we use the same energy range as in FERMI/LAT BRIGHT SOURCE LIST 5.2 12.6 5.1 12.4 5 1 s ] 12.8 12.2 53 4.9 2 Flux [10 ph cm 12 6 Flux [10 6ph cm 2 s 1] No. 1, 2009 11.8 11.6 11.4 4.7 4.6 4.5 4.4 11.2 11 4.8 4.3 220 230 240 250 260 270 280 290 300 Figure 5. Light curve of Vela (0FGL J0835.4−4510) with fluxes from a single power-law fit and purely statistical error bars. Each interval is approximately one week. The dashed line is the average value. Because Vela is very bright it would have been classified as variable using the statistical errors only, but the flux dispersion is only 2.3% beyond statistical. The gray area shows the 3% systematic error we have adopted. Note that because this analysis uses a power-law model over the entire range from 200 MeV to 100 GeV it grossly overestimates the true flux of Vela, but this effect does not depend on time. Section 3.3 (200 MeV to 100 GeV) to study variability. To avoid ending up with too large error bars in relatively short time intervals, we froze the spectral index of each source to the best fit over the full interval. Sources do vary in spectral shape as well as in flux, of course, but we do not aim to characterize source variability here, just detect it. It is very unlikely that a true variability in shape will be such that it will not show up in flux at all. We split the full three month interval into Nint = 12 intervals of a little more than one week. This preserves some statistical precision for the moderately bright sources we are dealing with here, while being sensitive in the right timescale for flaring blazars. Because we do not expect the diffuse emission to vary, we freeze the spectral adjustment of the Galactic diffuse component to the local (in the same RoI) best fit over the full interval. We need to leave the normalization of at least one diffuse component free (just to adapt to the natural Poisson variations of the background). Because it was not obvious which one to freeze we decided to leave both (Galactic and isotropic) free in each interval. So in the end the fitting procedure is the same as in Section 3.3 except that all spectral shape parameters are frozen. The faint sources were left free (even when not significant in the current interval) as well as the bright ones. As in Section 3.4, it often happens that a source is not significant in all intervals. To preserve the variability index (Equation (2)) we keep the best-fit value and its estimated error even when the source is not significant. This does not work, however, when the best fit is very close to zero because in that case the log(likelihood) as a function of flux is very asymmetric. Whenever TS < 1 we compute the 1σ upper limit and replace the error estimate with the difference between that upper limit and the best fit. This is an estimate of the error on the positive side only. The best fit itself is retained. Figures 5 and 6 show the fluxes derived for the Vela and Geminga pulsars as a function of time. As the brightest persistent sources, Vela and Geminga provide a reference for nonvariability. Based on these light curves, we estimate that the instrument 4.2 220 230 240 250 260 270 280 290 300 Figure 6. Same as Figure 5 for Geminga (0FGL J0634.0+1745). That pulsar’s flux dispersion is 1.6% beyond statistical. Its variability index (Equation (2)) would not have exceeded the threshold even with pure statistical errors. and processing (event classification) are stable on timescales of weeks to 2% relative precision. To be conservative we have added in quadrature a fraction frel = 3% of the average flux Fav to the error estimates (for each one week time interval) used to compute the variability index. Figure 7 shows the flux derived for the AO 0235+164 blazar as a function of time. In contrast to the steady pulsars, many of the blazars detected by the LAT show strong variability. The variability index is defined as a simple χ 2 criterion: V =  (Fi − Fav )2 , σi2 + (frel Fav )2 i (2) where i runs over the 12 intervals and σi is the statistical uncertainty in Fi . Since Fav is not known a priori, this parameter is expected, in the absence of variability, to follow a χ 2 distribution with 11 (= Nint − 1) degrees of freedom. We set the variability flag true whenever the probability of getting the value of V or more by chance is less than 1% (so that we expect two false positives over the sample of 205 sources). This corresponds to V > 24.7. This variability index is robust for the bright sources considered here, although for less significant sources methods that handle upper limits will be needed. Figure 8 shows the relative variability of the sources. It is defined from the excess variance on top of the statistical and systematic fluctuations:   2 2 i σi i (Fi − Fav ) 2 δF /F = − − frel . (3) 2 2 (Nint − 1)Fav Nint Fav The typical relative variability is 50%, with only a few strongly variable sources beyond δF /F = 1. The dotted lines show how the relative variability depends on the variability index as a function of TS, assuming that the relative precision on flux is the same for all sources (we get a relative precision of 0.16 for the flux on average at TS = 100). So this means that the criterion we use is not sensitive to relative variations smaller than 60% at TS = 100. That limit goes down to 20% as TS increases to 1000. Note that because the relative precision on flux is not exactly the same for all sources the cutoff, expressed in relative variability, 54 ABDO ET AL. Vol. 183 1 s ] 1 6 Flux [10 ph cm 2 0.8 0.6 0.4 0.2 0 220 230 240 250 260 270 280 290 300 Figure 7. Same as Figure 5 for AO 0235+164 (0FGL J0238.6+1636), a variable blazar. Note the difference in scale from the Vela and Geminga light curves. is not sharp. It is clear that we must be missing many variable AGNs below TS = 1000. Sixty-six bright sources (one third of the sample) are declared variable. This is far more than the two false positives expected, so most of them are truly variable. This level of variability is not particularly surprising as blazars are known to be strongly variable on timescales of days to weeks. We emphasize that sources not flagged may also show variability at lower amplitude or different timescales than used for this test. We refer to these other sources as “nonvariable” (on weekly timescales) rather than “steady.” 3.6. Limitations and Systematic Uncertainties A limitation of this work is that we did not attempt to test for source extension. All sources are assumed to be point-like. This is true for all known source populations in the GeV range (see Section 4). On the other hand the TeV instruments have detected many extended sources in the Galactic plane, mostly pulsar wind nebulae (PWNe) and supernova remnants (SNR; e.g., Funk 2005; Abdo et al. 2007). The current level of LAT exposure cannot address source extension at the level seen by the TeV telescopes. We have addressed the issue of systematics for localization in Section 3.2. This section deals with the systematic uncertainties on flux estimates. An obvious one is the power-law representation within each energy band. If one source had a very curved spectrum (like a spectral line) its flux estimate certainly would be inaccurate. Our experience with those sources for which more detailed studies have been made, though, is that the current estimates are fully acceptable. Beyond that, there are two main sources of systematics: the imperfect knowledge of the instrument so early into the mission, and the imperfect modeling of the diffuse emission. The fluxes are calculated using pre-launch calibrations (designated P6_V1) based on Monte Carlo simulations and a beam test at CERN (Atwood et al. 2009). In flight, the presence of pile-up signals in the LAT tracker and calorimeter left by earlier particles was revealed in periodic trigger events. This effect leads to a reduction of the actual acceptance as compared to the pre-launch prediction as fewer events pass the rejection cuts, most notably for photons below 300 MeV. The magnitude of this reduction is still under investigation, but the fluxes reported here may be lower than the true ones by as much as 35% below Figure 8. Relative source variability plotted as a function of the variability index (Equation (3)). The vertical dashed line shows where we set the variable source limit. The √ horizontal dashed line is the maximum relative variability that can be measured Nint − 1. The dotted lines show how the variability index depends on δF /F at our threshold (TS = 100) and for brighter sources (TS = 1000). At a given TS, the lower right part of the diagram is not accessible. For more details see Section 3.5. 1 GeV and 15% above 1 GeV. Because of the current uncertainty, no correction has been applied to the results; these effects are being assessed in detail, and will then be included in a reprocessing of the data. This uncertainty applies uniformly to all sources. Our relative errors (comparing one source to another or the same source as a function of time) are much smaller, as indicated in Section 3.5. It is interesting to note that the flux above 100 MeV that the LAT finds for the three historical pulsars (Vela, Geminga, and the Crab) is actually very close to that reported in the 3EG catalog (Hartman et al. 1999). Geminga and the Crab are within 1σ , and the LAT flux for Vela (9.15 ×10−6 photons cm−2 s−1 ) is 11% higher than that of EGRET. This implies that the bias may be on the low side of our estimate unless EGRET also underestimated the source flux. The diffuse emission is the other important source of uncertainties. Contrary to the former, it does not affect all sources equally. It is essentially negligible (i.e., smaller than the statistical errors) outside the Galactic plane (|b| > 10◦ ) where the diffuse emission is faint and varying on large angular scales. It is also not much of a worry in the high band (> 1 GeV) where the PSF is sharp enough that the bright sources dominate the background under the PSF. But it is a serious issue inside the Galactic plane (|b| < 10◦ ) in the low band (< 1 GeV) and particularly inside the Galactic ridge (|l| < 60◦ ) where the diffuse emission is strongest and very structured, following the molecular cloud distribution. It is not easy to assess precisely how large the uncertainty is, for lack of a proper reference. We have tried re-extracting the source fluxes assuming a very different diffuse model, and the results tend to show that the systematic uncertainty more or less follows the statistical one (i.e., it is larger for fainter sources) and is of the same order. We have not increased the errors accordingly, though, because this alternative model does not fit the data as well as the reference model where the differences in the source fluxes are largest. The net result of these considerations is that we expect our high-energy fluxes to be reasonably accurate, but the low-energy fluxes are not as reliable and should be treated with particular caution in the Galactic ridge. No. 1, 2009 FERMI/LAT BRIGHT SOURCE LIST 55 Table 1 Catalogs Used for Automatic Source Association Name Objects Ė/d 2 ) Pulsars (high Pulsars (low Ė/d 2 ) PWN SNR HMXB LMXB Microquasars Globular clusters Blazars (CGRABS) Blazars (BZCAT) Flat Spectrum Radio Sources (CRATES) 3EG catalog EGR AGL Selection Pprior Ė/d 2 5 1033 Reference cm−2 s−1 100 1527 69 265 114 187 15 147 0.29 0.044 0.5 0.033 0.17 0.19 0.5 0.5 1625 2686 10272 0.14 0.043 0.022 Healey et al. (2008) Massaro et al. (2009) Healey et al. (2007) n.a. n.a. n.a. Hartman et al. (1999) Casandjian & Grenier (2008) Pittori et al. (2009) 271 189 40 > erg Ė/d 2  5 1033 erg cm−2 s−1 Manchester et al. (2005) Manchester et al. (2005) Roberts (2005)a Green (2006)b Liu et al. (2006) Liu et al. (2007) Paredes (2006) Harris (1996) Notes. For clarity the table has been divided into Galactic, extragalactic, and γ -ray source catalogs. a http://www.physics.mcgill.ca/ pulsar/pwncat.html. b http://www.mrao.cam.ac.uk/surveys/snrs/. 4. SOURCE ASSOCIATION AND IDENTIFICATION Even with the superior angular resolution of LAT compared to previous generation γ -ray telescopes, the source location accuracy is not good enough to draw firm conclusions based on positional coincidence in most cases. A typical LAT error circle contains multiple stars, galaxies, X-ray sources, infrared sources, and radio sources. Determination of the nature of a given LAT source must therefore rely on more information than only simple location. The following two principles lead the search for counterparts. 1. Variability is a powerful diagnostic, particularly considering that many γ -ray sources are known to be variable. Searches for periodic variability (such as rotational and orbital motion) offer opportunities for unique identifications. Determining variability correlated with that seen at other wavelengths is another approach. 2. LAT γ -ray sources are necessarily nonthermal objects involving large energy transfers. Physical properties of any candidate counterpart must be consistent with generation of a significant luminosity of gamma radiation. In this analysis, the LAT team makes a clear distinction between a source identification and an association with an object at another wavelength. A firm identification of a source is based on a timing characteristic such as a periodicity for a pulsar or binary or a variability correlated with observations at another wavelength in the case of a blazar. An association is made for a statistically improbable positional coincidence of a plausible γ -ray-producing object with a LAT source. 4.1. Automated Source Associations In anticipating the large number of γ -ray sources that will be detected by the LAT in the course of the mission, we implemented an automated source association pipeline that attempts to make quantified associations between LAT sources and potential counterparts. In its implementation for the Bright Gamma-Ray Source List the pipeline is almost exclusively based on positional coincidence, yet is driven by past knowledge about GeV source classes (pulsars and blazars) and physical expectations (such as total luminosity and nonthermal emission implying particle acceleration). Future implementations will also include figureof-merit (FoM) approaches (Sowards-Emmerd et al. 2003) but these first require careful training on firmly identified source classes. For each LAT source the probability of association with a source in the counterpart catalog is estimated using a Bayesian approach (e.g., de Ruiter et al. 1977; Sutherland & Saunders 1992) that considers the spatial match between LAT source and counterpart in light of the position uncertainty r95 and the chance coincidence probability as inferred from the local source density in the counterpart catalog. Specifically, we calculate the posterior probability of association Ppost  −1 2 1 − Pprior πρ r95 ∆ = 1+ e , Pprior 2.996 (4) 2 where ρ is the local counterpart density, ∆ = 2.996 × r 2 /r95 , r is the angular separation between LAT source and catalog counterpart, and Pprior is the prior probability of the association that we use here as a constant tuning parameter whose value is adjusted for each counterpart catalog to give an approximately constant false association rate among the catalogs considered. Since the value of Ppost depends on the choice of Pprior we arbitrarily define a counterpart as a possible association if Ppost  0.5. We applied our pipeline to random realizations of plausible LAT catalogs64 in order to find for each counterpart catalog the value of Pprior that does not produce more than a single spurious association. Table 1 summarizes the catalogs that have been used in our automatic association procedure. We also quote the prior probabilities that have been employed and give the total number of objects in each catalog. Note that we make an exception to our procedure when we cross-correlate the EGRET 3EG and EGR and AGILE AGL catalogs with the LAT sources. Since in these cases the uncertainties in the localization of the counterparts is worse than for the LAT sources we consider all EGRET 64 The plausible LAT catalogs contained 1000 sources of which 75% were distributed isotropically over the sky and 25% were distributed along the Galactic plane following a two-dimensional Gaussian shaped density profile with σ = 40 deg in longitude and σ = 2 deg in latitude. For each source an error radius r95 of 0.2 deg has been assumed. 56 ABDO ET AL. 20 Number of objects 15 10 5 0 0 1 2 angular separation (sigma) 3 Figure 9. Distribution of angular separations between LAT sources and counterpart catalog associations expressed as σ = 0.405 r/r95 . The expected distribution in the case that all sources have been correctly associated is given as the dotted line. The peak at somewhat lower angular separation than the dotted prediction might indicate slightly better position determinations on average than (conservatively) assumed in this paper. Conversely, the expected distribution in the case that all sources are spurious associations is given as the dashed line. Vol. 183 previous section. This method uses a FoM approach similar to the one described by Healey et al. (2008), based not only on positional proximity but also on radio spectral index, X-ray flux, and radio flux (Abdo et al. 2009c). Details of this association procedure, including the calculated probabilities from both the FoM and automated association approaches, can be found in that paper. Although most of the associations are found by both methods, about 11% are found only by one of the two. In order to maintain consistency with the LAT AGN paper (Abdo et al. 2009c), we show any association found by either method. It should be emphasized, however, that Abdo et al. (2009c) chose to apply the AGN analysis only to parts of the sky with Galactic latitudes more than 10◦ from the plane in order to have a more uniform sample, while the present analysis covers the entire sky. AGN are seen by LAT at lower Galactic latitudes, because the Galaxy is largely transparent to γ -rays. Due to Galactic extinction and source confusion, AGN identification is more difficult at low latitudes. Some of the unassociated LAT sources in this part of the sky can be expected to have AGN counterparts in further analysis, which is beyond the scope of this paper. 4.2.2. Firm Identifications and AGILE sources as possible counterparts if the LAT and counterpart separation is less than the quadratic sum of their 95% confidence error radii. The pulsar catalog (the ATNF Pulsar Catalog; Manchester et al. 2005) is special in that we split it into high and low Ė/d 2 subsamples, where Ė is the rate of energy loss of the pulsar and d is the distance. High Ė/d 2 has been proposed as a good estimator of a pulsar’s γ -ray visibility (Smith et al. 2008), and downselecting the catalog to the 100 best candidates allows for a relatively large prior probability without inflating the number of false positives. For the remaining pulsars our Monte Carlo simulations required a much smaller prior probability (to keep the chance coincidences low) at the expense of reducing the number of potential associations. This procedure can be considered as a simple binary FoM approach which favors revealing high Ė/d 2 counterparts of LAT sources. The performance of our association scheme is illustrated in Figure 9, which shows the distribution of normalized angular separations between LAT sources and counterparts; the normalization is done with respect to the measured localization uncertainty (Section 3.2). We also show the expected distribution for the case that all sources have been correctly associated (dotted line) or are spurious associations (dashed line). Obviously, the observed distribution clearly follows the first trend, suggesting that most of our associations are indeed reasonable and that our efforts to reduce the number of false positives were successful. We note that the histogram shows a slight trend to smaller angular separations than expected, which might result from a slight overestimation of our source localization uncertainties. 4.2. Alternate Associations, Firm Identifications, and Special Cases 4.2.1. Active Galactic Nuclei AGNs have been recognized since the EGRET era as a well defined class of gamma-ray sources. For this reason, we have adopted an alternate method of finding AGN associations beyond the automated association procedure described in the For this early source list from the LAT, we have taken the conservative view that association, even with high probability, is not equivalent to firm identification. Error circles are still large compared to source localization at longer wavelengths. We adopt the approach that firm identification for the 0FGL sources is limited to those for which variability can unambiguously establish the source. Firm identifications of pulsars are based on seeing the pulsations in the γ -ray data with high confidence. Using several statistical tests, we require that the γ -ray distribution in pulsar phase be inconsistent with random at a probability level of 10−6 or smaller. Examples are the six pulsars confirmed from the EGRET era, the radio-quiet pulsar found in the CTA 1 SNR (Abdo et al. 2008), PSR J0030+0451 (Abdo et al. 2009o), PSR J1028−5819 (Abdo et al. 2009f), and PSR J2021+3651 (Abdo et al. 2009n). In total the 0FGL source list includes 30 firm pulsar identifications. One third of the sources within 10◦ of the Galactic plane have now been identified with pulsars. The high-mass X-ray binary (HMXB) system LSI +61 303 is firmly identified based on the observation of the orbital period of the binary system (Abdo et al. 2009j). A search for periodicity in the similar source LS5039 is still in progress. Firm identifications of AGNs depend on finding correlated multiwavelength activity. This work is ongoing. 4.2.3. Special Cases – Pulsar Wind Nebulae and Supernova Remnants SNRs and PWNe that are positionally correlated with LAT sources are not listed as individual associations in the main source list table. Statistical indications are that SNRs that were coincident with EGRET sources are significantly correlated with the 0FGL sources. However, the large number of pulsars detected by the LAT, including radio-quiet pulsars (e.g., Abdo et al. 2008), suggests that even a positional coincidence with an SNR of an age, distance, and environment plausible for a γ -ray source may be due to a γ -ray pulsar. Of the 0FGL sources positionally associated with PWNe or SNRs, approximately 40% have already been found to contain γ -ray pulsars. At the present level of sensitivity for the LAT-detected pulsars, only the Crab has shown evidence for off-pulse emission that can No. 1, 2009 FERMI/LAT BRIGHT SOURCE LIST 57 Table 2 Potential Associations for Sources Near SNRs and PWNe Name 0FGL J0617.4+2234 J1018.2−5858 J1106.4−6055 J1615.6−5049 J1648.1−4606 J1714.7−3827 J1801.6−2327 J1814.3−1739 J1834.4−0841 J1855.9+0126 J1911.0+0905 J1923.0+1411 J1954.4+2838 l b Association 189.08 284.30 290.52 332.35 339.47 348.52 6.54 13.05 23.27 34.72 43.25 49.13 65.30 3.07 −1.76 −0.60 −0.01 −0.71 0.10 −0.31 −0.09 −0.22 −0.35 −0.18 −0.40 0.38 SNR G189.1+3.0 (IC 443) SNR G284.3−1.8 (MSH 10-53), PSR J1013−5915 SNR G290.1−0.8 (MSH 11-61A), PSR J1105−6107 SNR G332.4+0.1, PWN G332.5−0.28, PSR B1610−50 PSR J1648−4611 SNR G348.5+0.1 SNR G6.4−0.1 (W28) PWN G12.82−0.02 SNR G23.3−0.3 (W41) SNR G34.7−0.4 (W44) SNR G43.3−0.2 SNR G49.2−0.7 (W51) SNR G65.1+0.6 Notes. See the text, Section 4.2.3. These sources are marked with a † in the source list table. They may be pulsars rather than the SNR or PWN named. Table 3 LAT Bright Source List Description Column Description Name 0FGL JHHMM.m+DDMM, constructed according to IAU Specifications for Nomenclature; m is decimal minutes of R.A.; in the name R.A. and decl. are truncated at 0.1 decimal minutes and 1′ , respectively Right Ascension, J2000, deg, 3 decimal places Declination, J2000, deg, 3 decimal places Galactic Longitude, deg, 3 decimal places Galactic Latitude, deg, 3 decimal places Radius of 95% confidence region, deg, 3 decimal places Square root of likelihood TS from 200 MeV–100 GeV analysis, used for the TS > 100 cut, 1 decimal place Flux 100 MeV to 1 GeV (i.e., log10 E = 2–3), 10−8 cm−2 s−1 , 2 decimal places 1σ uncertainty on F23 , same units and precision. A 0 in this column indicates that the entry in the F23 flux column is an upper limit. Square root of TS for the 100 MeV to 1 GeV range, 1 decimal place Flux for 1 GeV to 100 GeV (i.e., log10 E = 3–5), 10−8 cm−2 s−1 , 2 decimal places 1σ uncertainty on F35 , same units and precision Square root of TS for the 1 GeV to 100 GeV range, 1 decimal place T indicates < 1% chance of being a steady source on a weekly timescale; see Section 3.5 Identification or positional associations with 3EG, EGR, or AGILE sources Like “ID” in 3EG catalog, but with more detail (see Table 4). Capital letters indicate firm identifications; lower-case letters indicate associations. Identification or positional associations with potential counterparts Reference to associated paper(s), R.A. Decl. l b θ95 TS1/2 F23 ∆F23 1/2 TS23 F35 ∆F35 1/2 TS35 Var. γ -ray Assoc. Class ID or Assoc. Ref. be attributed to a PWN or SNR. Until the possibility of pulsed emission for such sources can be ruled out, we are reluctant to make any claims about individual PWNe or SNRs as possible LAT detections. Effectively, the high rate of pulsar detections increases the burden of proof for PWN and SNR candidates, for example, via studies of source extents. Table 2 shows the 0FGL sources that are associated positionally with PWNe and SNRs, plus four pulsars that do not (yet) show evidence of γ -ray pulsation. Torres et al. (2003) considered several of the SNRs in Table 2 in terms of their potential to be γ -ray counterparts to unidentified low-latitude EGRET sources. In the case of SNR G284.3−1.8 they argued that PSR J1013−5915 was more probably the γ -ray source. 5. THE SOURCE LIST Table 3 is a description of the columns in the source list table. Within the table, sources that have firm identifications or tentative associations are listed by class. Table 4 describes those classes. The bright source list itself is presented as a single table (Table 5). Figure 10 shows the locations of the 205 bright sources in Galactic coordinates. All associations with specific source classes are also shown. Figure 11 is an enlargement of the Table 4 LAT Bright Source List Source Classes Class Description PSR pwn hxb bzb bzq bzu rdg glb † Pulsar Pulsar wind nebula High-mass X-ray binary (black hole or neutron star) BL Lac type of blazar FSRQ type of blazar Uncertain type of blazar Radio galaxy Globular cluster Special case—potential association with SNR or PWN (see Table 2) Notes. Designations shown in capital letters are firm identifications; lower-case letters indicate associations. In the case of AGNs, many of the associations have high confidence (Abdo et al. 2009c). Among the pulsars, those with names beginning with LAT are newly discovered by the LAT. bright source map, showing the region of the inner Galaxy. This list is available as a FITS file from the Fermi Science Support Center. 58 ABDO ET AL. Vol. 183 +90 +180 180 o variable x not variable 90 Figure 12. Locations of variable (circles) and nonvariable (crosses) 0FGL sources, using the definition of variability in Section 3.5. The analysis is sensitive to variations on timescales of weeks to ∼two months. Figure 10. The LAT bright source list, showing the locations on the sky (Galactic coordinates in Aitoff projection) coded according to the legend. Although quantitative spectral information is not presented, the colors of the symbols indicate relative spectral hardness on a sliding scale. Symbols more blue in color indicate sources with harder spectra than those that are more red. Galactic Latitude [deg] 30 15 0 15 30 90 75 60 45 30 15 0 345 330 315 300 285 270 Galactic Longitude [deg] Figure 11. The LAT bright source list, showing the locations of sources in the inner Galaxy. The legend is the same as in Figure 10. 6. DISCUSSION As is clear from the references in this paper, much of the work on the early data from Fermi/LAT is still in progress. In particular, we re-emphasize several caveats for use of this bright source list. 1. Ongoing efforts to understand the calibration and improve the analysis techniques are underway. In many respects, therefore, the 0FGL source list is quite preliminary. Significant improvements are expected before the construction of the first full LAT catalog. 2. The GALPROP diffuse model used in the analysis is still evolving. Matching the model to the large-scale emission is an iterative process. The diffuse model is particularly important for sources near the Galactic plane. 3. This source list is limited to high-confidence detections. It is not a full catalog. 4. The 0FGL list information in two broad energy bands is not appropriate for detailed spectral modeling. 5. This work is a “snapshot” of the LAT results covering only the observation time period 2008 August to October. 6. The use of the Diffuse class of events means that LAT has little sensitivity below 200 MeV for this particular analysis. As noted by Abdo et al. (2009c) in their analysis of AGNs, the LAT is more sensitive to hard-spectrum sources than previous satellite instruments. Despite these issues, the present work demonstrates the power of the LAT to make high-energy γ -ray observations and shows its potential for future discoveries. Although we feel it premature to draw far-reaching conclusions, some results stand out. 6.1. Characteristics of the 0FGL Sources 1. Both Galactic and extragalactic populations are visible. Seventy-three sources are found within 10◦ of the Galactic plane, where they exhibit a characteristic concentration in the inner Galaxy; 132 are seen at higher Galactic latitudes. 2. Sixty-six of the bright LAT sources show solid evidence of variability on weekly timescales during this three month interval. Figure 12 shows the locations of variable and nonvariable sources in Galactic coordinates. 3. The typical error radii for the sources (95% confidence) are less than 10′ . 4. The Galactic latitude distribution of unassociated/ unidentified γ -ray sources is very narrow (FWHM <0.◦ 5). If we assume a scale height for a Galactic population of 40 pc (Guibert et al. 1978), such a narrow latitude distribution points to a Galactic γ -ray source population with average distance in excess of 40/sin 0.◦ 5, namely, 4.5 kpc. 6.2. Comparisons with Other High-Energy γ -Ray Results Before Fermi, the EGRET results represented the most complete view of the high-energy sky, but those results applied to the 1991–2000 era. In light of the variability seen in the EGRET γ -ray sources, significant differences were expected. A contemporaneous mission to Fermi is AGILE, which began operations over a year before Fermi and continues to operate. Here is a summary of comparisons with these missions: 1. Of the 205 0FGL sources, 60 have nearby counterparts (the LAT source 95% uncertainty overlapping that of the EGRET source) found by the automated analysis in the 3EG catalog (271 sources); 43 in the EGR catalog (189 sources). Most of the sources seen by EGRET in the 1990s were not seen by LAT as bright sources in 2008. Approximately 40% of the bright LAT sources off the plane that have no former EGRET counterparts are found to be variable. 2. EGRET found few sources with flux less than 10 × 10−8 photons (E > 100 MeV) cm−2 s−1 . A number of the 0FGL sources have fluxes well below this value (e.g., 0FGL J0033.6−1921). Such sources would not have been visible to EGRET. 3. Some sources, such as 0FGL J0428.7−3755, associated with blazar PKS 0426−380, have flux values well above the EGRET threshold but were not seen by EGRET and yet are not noted as being variable in the 0FGL data. Such sources serve as a reminder that blazars are variable on many timescales, and the 0FGL sample covers only three months. No. 1, 2009 FERMI/LAT BRIGHT SOURCE LIST 59 Table 5 LAT Bright Source List Name 0FGL R.A. J0007.4+7303 1.852 Decl. l b θ95 100 MeV to 1 1 GeV to 100 GeV GeV √ √ √ TS F23 ∆F23 TS23 F35 ∆F35 TS35 Var. 73.065 119.690 10.471 0.054 64.6 32.4 1.3 36.2 6.14 0.27 55.9 J0017.4−0503 4.358 −5.054 101.273 −66.485 0.252 14.7 11.8 1.4 J0025.1−7202 6.295 −72.042 305.786 −44.940 0.163 12.7 5.1 2.0 15.3 0.27 0.07 6.8 0.56 0.10 7.1 11.5 J0030.3+0450 7.600 4.848 113.111 −57.622 0.138 18.7 8.9 J0033.6−1921 8.401 −19.360 94.215 −81.220 0.147 10.7 1.4 0.0 0.0 13.6 0.71 0.10 4.6 0.36 0.07 14.9 10.1 J0036.7+5951 9.177 59.854 121.081 −2.965 0.144 10.3 6.7 3.1 5.6 0.48 0.10 8.6 J0050.5−0928 12.637 −9.470 122.209 −72.341 0.130 20.5 8.1 1.3 15.6 0.72 0.10 14.9 J0051.1−0647 12.796 −6.794 122.751 −69.666 0.127 15.7 6.6 1.4 10.4 0.57 0.09 12.0 J0100.2+0750 15.051 7.844 126.716 −54.963 0.110 11.1 2.5 J0112.1+2247 18.034 22.789 129.148 −39.832 0.134 17.6 5.4 0.0 0.7 2.8 0.31 0.07 10.8 0.65 0.09 10.4 14.4 J0118.7−2139 19.676 −21.656 172.990 −81.728 0.164 17.8 7.0 1.1 14.5 0.52 0.09 12.2 J0120.5−2703 20.128 −27.056 213.951 −83.529 0.140 11.8 2.3 0.8 6.6 0.33 0.07 10.3 J0136.6+3903 24.163 39.066 132.446 −22.969 0.087 12.5 5.9 0.0 3.6 0.45 0.08 12.3 J0137.1+4751 24.285 47.854 130.818 −14.317 0.120 18.8 10.0 1.6 12.3 0.78 0.10 15.4 J0144.5+2709 26.142 27.159 137.248 −34.231 0.209 10.4 1.7 0.5 6.6 0.32 0.07 7.4 J0145.1−2728 26.289 −27.478 217.694 −78.067 0.243 13.4 9.2 1.3 13.7 0.26 0.07 6.9 J0204.8−1704 31.219 −17.068 186.072 −70.274 0.163 16.6 10.2 1.3 15.0 0.44 0.08 10.8 J0210.8−5100 32.706 −51.013 276.083 −61.776 0.070 34.1 21.4 1.2 28.2 1.35 0.14 22.2 J0217.8+0146 34.467 1.2 16.1 0.82 0.11 16.7 J0220.9+3607 35.243 36.121 142.504 −23.325 0.225 12.3 10.7 1.3 13.1 0.22 0.06 6.0 J0222.6+4302 35.653 43.043 140.132 −16.763 0.054 47.4 24.0 1.4 32.0 2.61 0.18 37.4 J0229.5−3640 37.375 −36.681 243.801 −67.189 0.138 19.2 13.7 1.5 16.9 0.45 0.08 10.9 J0238.4+2855 39.600 28.923 149.521 −28.368 0.193 10.9 8.3 9.3 0.34 0.08 7.5 J0238.6+1636 39.663 16.613 156.775 −39.112 0.052 85.7 60.7 2.1 64.3 6.81 0.29 62.5 J0240.3+6113 40.093 61.225 135.661 37.4 3.34 0.23 27.6 1.768 162.139 −54.389 0.106 21.7 8.9 1.075 1.6 0.069 42.3 70.3 2.5 J0245.6−4656 41.423 −46.934 262.019 −60.098 0.192 11.4 5.3 0.8 9.0 0.32 0.07 8.1 J0303.7−2410 45.940 −24.176 214.764 −60.119 0.174 12.3 2.5 0.9 7.8 0.38 0.08 10.2 J0320.0+4131 50.000 41.524 150.601 −13.230 0.086 29.7 16.6 1.4 21.6 1.60 0.15 22.6 J0334.1−4006 53.546 −40.107 244.710 −54.088 0.152 13.2 4.5 9.0 0.39 0.08 10.7 J0349.8−2102 57.465 −21.046 214.385 −49.035 0.157 21.2 16.7 1.6 20.4 0.56 0.09 10.9 J0357.5+3205 59.388 32.084 162.712 −16.056 0.147 14.9 10.4 1.8 J0407.6−3829 61.923 −38.491 241.360 −47.751 0.142 13.5 6.5 1.3 13.6 0.64 0.10 11.2 0.41 0.08 10.5 9.6 J0412.9−5341 63.230 −53.686 263.001 −44.716 0.206 10.7 5.7 8.8 7.9 1.4 1.3 0.29 0.07 γ -Ray Assoc. Class ID or Assoc. · · · 3EG J0010+7309 PSR LAT PSR J0007+7303 EGR J0008+7308 1AGL J0006+7311 T ··· bzq CGRaBS J0017−0512 ··· ··· glb NGC 104 47 Tuc ··· ··· PSR PSR J0030+0451 ··· ··· bzb BZB J0033−1921 KUV 00311−1938 ··· ··· bzb BZB J0035+5950 1ES 0033+595 T ··· bzb CGRaBS J0050−0929 PKS 0048−097 T ··· bzq CGRaBS J0051−0650 PKS 0048−071 ··· ··· bzu CRATES J0100+0745 ··· ··· bzb CGRaBS J0112+2244 S2 0109+22 T ··· bzq CGRaBS J0118−2141 PKS 0116−219 ··· ··· bzb CGRaBS J0120−2701 PKS 0118−272 ··· ··· bzb BZB J0136+3905 B3 0133+388 T ··· bzq CGRaBS J0136+4751 DA 55 ··· ··· bzb CRATES J0144+2705 TXS 0141+268 T ··· bzq CGRaBS J0145−2733 PKS 0142−278 ··· ··· bzq CGRaBS J0204−1701 PKS 0202−17 T 3EG J0210−5055 bzq CGRaBS J0210−5101 EGR J0210−5058 PKS 0208−512 T ··· bzq CGRaBS J0217+0144 PKS 0215+015 ··· ··· bzq CGRaBS J0221+3556 B2 0218+35 T 3EG J0222+4253 bzb BZB J0222+4302 EGR J0223+4300 3C 66A T ··· bzq BZQ J0229−3643 PKS 0227−369 ··· ··· bzq CGRaBS J0237+2848 B2 0234+28 T 3EG J0237+1635 bzb CGRaBS J0238+1636 AO 0235+164 T EGR J0240+6112 HXB LS I+61 303 1AGL J0242+6111 ··· ··· bzu CRATES J0246−4651 PKS 0244−470 ··· ··· bzb CRATES J0303−2407 PKS 0301−243 T ··· rdg CGRaBS J0319+4130 NGC 1275 ··· ··· bzb CGRaBS J0334−4008 PKS 0332−403 ··· ··· bzq CGRaBS J0349−2102 PKS 0347−211 ··· ··· PSR LAT PSR J0357+32 T ··· bzq CRATES J0406−3826 PKS 0405−385 ··· ··· bzu CRATES J0413−5332 Ref. 1 ··· 2 3, 4 ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· 5 ··· ··· 6 ··· ··· 7 ··· ··· 60 ABDO ET AL. Vol. 183 Table 5 (Continued) Name 0FGL R.A. Decl. l b θ95 100 MeV to 1 GeV 1 GeV to 100 GeV √ √ √ TS F23 ∆F23 TS23 F35 ∆F35 TS35 Var. J0423.1−0112 65.785 −1.204 195.131 −33.092 0.143 11.5 10.1 3.1 10.3 J0428.7−3755 67.193 −37.923 240.689 −43.597 0.079 39.6 21.1 1.6 28.4 J0449.7−4348 72.435 −43.815 248.780 −39.859 0.082 28.4 10.9 1.4 16.9 J0457.1−2325 74.288 −23.432 223.739 −34.880 0.065 52.3 34.0 1.8 43.0 J0507.9+6739 76.985 67.650 143.772 15.905 0.058 13.2 0.0 4.8 J0516.2−6200 79.063 −62.000 271.376 −34.834 0.181 11.2 10.4 0.0 6.7 J0531.0+1331 82.761 13.528 191.385 −10.992 0.133 17.3 22.9 2.1 16.5 5.4 J0534.6+2201 83.653 22.022 184.562 −5.764 0.046 139.2 204.0 9.2 116.6 J0538.4−6856 84.612 −68.940 279.281 −31.713 0.325 17.8 19.1 2.6 J0538.8−4403 84.725 −44.062 250.057 −31.075 0.072 48.6 31.7 1.8 18.4 38.8 J0613.9−0202 93.485 −2.047 210.468 −9.274 0.175 10.0 J0614.3−3330 93.577 −33.500 240.513 −21.801 0.083 28.9 2.3 0.2 6.2 15.7 J0617.4+2234 94.356 22.568 189.079 3.066 0.063 50.7 43.1 2.5 36.2 6.8 5.4 J0631.8+1034 97.955 10.570 201.302 0.507 0.151 10.4 6.9 3.0 6.4 J0633.5+0634 98.387 6.578 205.041 −0.957 0.105 23.4 18.9 2.5 17.4 J0634.0+1745 98.503 17.760 195.155 4.285 0.043 283.2 286.2 3.8 206.0 J0643.2+0858 100.823 8.983 204.010 2.290 0.121 15.7 22.2 2.8 J0654.3+4513 103.590 45.220 171.228 19.369 0.075 29.2 19.1 1.6 15.8 22.5 J0654.3+5042 103.592 50.711 165.676 21.107 0.083 15.6 J0700.0−6611 105.016 −66.199 276.778 −23.809 0.182 10.1 4.3 3.5 1.2 0.0 8.0 5.6 J0712.9+5034 108.231 50.575 166.688 23.900 0.146 11.2 J0714.2+1934 108.552 19.574 197.685 13.648 0.128 15.0 J0719.4+3302 109.869 33.037 185.139 19.855 0.141 12.3 3.0 9.5 7.1 0.7 1.6 1.5 6.1 12.0 9.8 J0722.0+7120 110.508 71.348 143.976 28.029 0.080 34.4 15.5 1.6 22.9 J0730.4−1142 112.607 −11.707 227.799 3.154 0.082 28.9 26.2 2.2 21.6 J0738.2+1738 114.575 17.634 201.933 18.081 0.137 11.9 3.3 1.4 8.2 J0818.3+4222 124.579 42.367 178.244 33.409 0.083 20.9 6.2 1.1 12.4 J0824.9+5551 126.239 55.859 161.981 35.142 0.214 10.6 10.7 1.3 11.5 J0826.0−2228 126.500 −22.480 243.964 8.941 0.144 12.7 8.2 5.1 1.3 J0835.4−4510 128.865 −45.170 263.560 −2.767 0.042 374.2 803.1 5.7 295.9 J0855.4+2009 133.857 20.162 206.810 35.974 0.178 15.1 7.4 1.3 12.3 0.38 0.08 8.8 γ -Ray Assoc. Class ID or Assoc. · · · 3EG J0422−0102 bzq CGRaBS J0423−0120 PKS 0420−014 1.99 0.17 29.6 · · · ··· bzb CGRaBS J0428−3756 PKS 0426−380 1.25 0.13 23.8 · · · ··· bzb CRATES J0449−4350 PKS 0447−439 2.61 0.19 34.5 T 3EG J0456−2338 bzq CGRaBS J0457−2324 EGR J0456−2334 PKS 0454−234 0.27 0.06 12.5 · · · ··· bzb BZB J0507+6737 1ES 0502+675 0.42 0.08 10.0 · · · 3EG J0512−6150 bzu CGRaBS J0516−6207 PKS 0516−621 0.76 0.12 10.6 T 3EG J0530+1323 bzq CGRaBS J0530+1331 EGR J0530+1331 PKS 0528+134 15.40 0.44 92.9 · · · 3EG J0534+2200 PSR PSR J0534+2200 EGR J0534+2159 pwn Crab 1AGL J0535+2205 0.64 0.12 9.5 · · · 3EG J0533−6916 · · · LMC 2.51 0.19 32.5 T 3EG J0540−4402 bzb CRATES J0538−4405 EGR J0540−4358 PKS 0537−441 1AGL J0538−4424 0.47 0.09 8.5 · · · ··· PSR PSR J0613−0200 1.64 0.15 25.4 · · · 3EG J0616−3310 · · · ··· EGR J0615−3308 4.99 0.27 38.8 · · · 3EG J0617+2238 † ··· EGR J0617+2238 1AGL J0617+2236 0.74 0.13 8.8 · · · ··· PSR PSR J0631+1036 1.60 0.17 16.9 · · · EGR J0633+0646 PSR LAT PSR J0633+0632 61.61 0.86 207.7 · · · 3EG J0633+1751 PSR PSR J0633+1746 EGR J0633+1750 Geminga 1AGL J0634+1748 0.84 0.14 9.9 T ··· ··· ··· 1.26 0.13 20.4 T ··· bzq CGRaBS J0654+4514 B3 0650+453 0.59 0.09 13.9 T ··· bzu CGRaBS J0654+5042 0.44 0.09 8.6 · · · ··· bzu CRATES J0700−6610 PKS 0700−661 0.36 0.08 9.6 T ··· bzb CGRaBS J0712+5033 0.51 0.09 10.6 T ··· bzq CLASS J0713+1935 0.37 0.08 8.9 T ··· bzq CRATES J0719+3307 TXS 0716+332 1.49 0.13 27.6 T 3EG J0721+7120 bzb CGRaBS J0721+7120 EGR J0723+7134 S5 0716+71 1AGL J0722+7125 1.74 0.16 21.7 T ··· bzq BZQ J0730−1141 PKS 0727−11 0.34 0.08 9.3 · · · 3EG J0737+1721 bzb CGRaBS J0738+1742 EGR J0737+1720 PKS 0735+178 0.79 0.11 16.7 · · · ··· bzb CGRaBS J0818+4222 OJ 425 0.17 0.05 5.2 T ··· bzq CGRaBS J0824+5552 TXS 0820+560 0.50 0.09 10.6 · · · ··· bzb BZB J0826−2230 PKS 0823−223 112.08 1.23 255.6 · · · 3EG J0834−4511 PSR PSR J0835−4510 EGR J0834−4512 Vela 1AGL J0835−4509 0.43 0.08 10.0 · · · 3EG J0853+1941 bzb CGRaBS J0854+2006 Ref. ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· 4 ··· ··· ··· 7 ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· 8 ··· No. 1, 2009 FERMI/LAT BRIGHT SOURCE LIST 61 Table 5 (Continued) Name 0FGL R.A. J0909.7+0145 137.446 Decl. 1.757 l b θ95 100 MeV to 1 1 GeV to 100 GeV GeV √ √ √ TS F23 ∆F23 TS23 F35 ∆F35 TS35 Var. Class 0.3 12.5 0.19 0.06 5.1 ··· ··· bzb J0910.2−5044 137.568 −50.743 271.569 −1.856 0.161 13.2 29.4 3.3 J0921.2+4437 140.320 44.617 175.809 44.876 0.128 15.2 8.8 1.4 12.3 1.10 0.18 12.5 0.42 0.08 8.7 10.8 T ··· ··· ··· ··· bzq J0948.3+0019 147.077 ··· bzq 236.530 38.549 0.287 12.8 8.9 1.4 12.6 0.24 0.07 5.9 T J0957.6+5522 149.424 55.375 158.605 47.939 0.092 24.0 8.4 1.3 17.0 0.79 0.10 18.8 · · · EGR J0957+5513 bzq J1012.9+2435 153.241 24.598 207.897 54.406 0.175 12.4 4.2 J1015.2+4927 153.809 49.463 165.473 52.727 0.062 23.8 7.9 0.9 1.5 9.2 0.35 0.08 10.7 1.00 0.11 8.6 22.4 T ··· ··· ··· bzq bzb 236.457 47.036 0.104 20.6 12.4 1.5 17.7 0.67 0.10 14.3 T ··· bzq 284.298 284.346 285.074 147.765 16.9 10.8 9.1 12.8 0.27 0.29 0.28 0.07 17.7 13.2 15.3 9.4 · · · 3EG J1013−5915 † ··· ··· ··· · · · 3EG J1027−5817 PSR ··· ··· bzq J1045.6−5937 161.409 −59.631 287.637 −0.548 0.123 19.5 30.7 4.6 J1047.6−5834 161.922 −58.577 287.385 0.508 0.138 18.5 43.9 0.0 15.5 2.29 0.25 13.4 2.53 0.25 14.8 16.8 J1053.7+4926 163.442 49.449 160.309 58.263 0.124 10.1 4.0 0.0 0.8 0.21 0.05 10.2 · · · 1AGL J1043−5931 · · · · · · 3EG J1048−5840 PSR EGR J1048−5839 ··· ··· bzb J1054.5+2212 163.626 22.215 216.968 63.049 0.178 11.2 4.0 J1057.8+0138 164.451 1.643 251.219 52.709 0.194 10.3 8.8 1.0 1.7 7.8 8.4 0.29 0.07 0.35 0.08 8.9 8.3 ··· ··· 31.2 4.07 0.25 33.8 ··· 8.2 0.33 0.07 10.0 ··· J1100.2−8000 165.057 −80.012 298.047 −18.212 0.285 12.1 10.8 2.2 10.3 0.31 0.08 6.3 T J1104.5+3811 166.137 38.187 179.868 65.056 0.055 47.1 13.3 1.3 23.9 2.61 0.17 40.9 ··· J1106.4−6055 166.605 −60.918 290.516 −0.604 0.251 10.8 16.9 5.2 8.6 1.43 0.22 9.0 ··· J1115.8−6108 168.967 −61.147 291.661 −0.384 0.214 12.1 25.7 5.3 J1123.0−6416 170.762 −64.268 293.519 −3.024 0.125 10.2 22.4 3.9 J1129.8−1443 172.454 −14.727 275.133 43.694 0.246 10.5 10.6 1.6 11.1 1.53 0.23 8.8 0.39 0.13 11.7 0.20 0.06 9.4 5.1 5.7 ··· T ··· J1146.7−3808 176.689 −38.149 289.170 22.988 0.185 10.4 3.1 1.2 6.5 0.38 0.08 8.2 ··· J1159.2+2912 179.800 29.216 199.605 78.307 0.192 14.6 9.1 1.0 13.9 0.29 0.07 7.7 ··· J1218.0+3006 184.517 30.108 188.826 82.097 0.099 27.4 9.0 1.0 14.9 1.41 0.14 24.1 T J1221.7+2814 185.439 28.243 201.593 83.336 0.101 24.0 6.5 0.8 13.1 1.03 0.12 20.6 T 289.975 64.355 0.083 52.0 63.9 2.6 54.2 1.61 0.14 25.7 T J1015.9+0515 153.991 J1018.2−5858 J1024.0−5754 J1028.6−5817 J1034.0+6051 154.564 156.001 157.166 158.504 0.317 228.640 31.262 0.273 11.6 9.3 γ -Ray Assoc. 5.254 −58.978 −57.903 −58.292 60.853 J1058.1−5225 164.527 −52.432 285.995 −1.765 −0.453 −0.459 49.122 6.673 0.113 0.106 0.079 0.209 22.4 13.9 16.0 14.8 36.7 42.0 27.6 7.1 0.073 43.7 25.6 2.0 J1058.9+5629 164.731 56.488 149.521 54.442 0.083 12.0 4.7 J1229.1+0202 187.287 2.045 4.8 0.0 0.0 1.3 1.8 3.06 2.55 2.56 0.33 J1231.5−1410 187.875 −14.179 295.642 48.410 0.087 30.9 9.3 J1246.6−2544 191.655 −25.734 301.571 37.125 0.168 11.7 7.2 1.2 1.4 18.8 1.71 0.15 9.1 0.37 0.08 25.6 8.8 ··· ··· J1248.7+5811 192.189 58.191 123.617 58.934 0.122 14.3 7.0 J1253.4+5300 193.369 53.001 122.229 64.125 0.154 12.1 5.2 1.6 1.5 9.9 8.1 0.39 0.07 0.33 0.07 12.2 9.6 ··· ··· 31.8 1.44 0.14 22.6 T J1256.1−0547 194.034 −5.800 305.081 57.052 0.079 36.8 28.3 1.8 ID or Assoc. OJ 287 CGRaBS J0909+0200 PKS 0907+022 ··· CGRaBS J0920+4441 RGB J0920+446 CGRaBS J0948+0022 PMN J0948+0022 CRATES J0957+5522 4C +55.17 CRATES J1012+2439 CGRaBS J1015+4926 1ES 1011+496 CRATES J1016+0513 PMN J1016+0512 ··· ··· PSR J1028−5819 CGRaBS J1033+6051 S4 1030+61 ··· PSR J1048−5832 BZB J1053+4929 MS 1050.7+4946 ··· bzb CLASS J1054+2210 ··· bzq CGRaBS J1058+0133 PKS 1055+018 3EG J1058−5234 PSR PSR J1057−5226 EGR J1058−5221 1AGL J1058−5239 ··· bzb CGRaBS J1058+5628 RXS J10586+5628 ··· bzb CGRaBS J1058−8003 PKS 1057−79 3EG J1104+3809 bzb CGRaBS J1104+3812 EGR J1104+3813 Mrk 421 1AGL J1104+3754 3EG J1102−6103 † ··· 1AGL J1108−6103 ··· ··· ··· ··· ··· ··· ··· bzq CRATES J1130−1449 PKS 1127−14 ··· bzq CGRaBS J1147−3812 PKS 1144−379 3EG J1200+2847 bzq CGRaBS J1159+2914 4C 29.45 ··· bzb CGRaBS J1217+3007 B2 1215+30 1AGL J1222+2851 bzb CGRaBS J1221+2813 W Com 3EG J1229+0210 bzq CGRaBS J1229+0203 EGR J1229+0203 3C 273 1AGL J1228+0142 EGR J1231−1412 · · · ··· ··· bzq CGRaBS J1246−2547 PKS 1244−255 ··· bzb PG 1246+586 ··· bzb CRATES J1253+5301 S4 1250+53 3EG J1255−0549 bzq CGRaBS J1256−0547 Ref. ··· ··· ··· 9 ··· ··· ··· ··· ··· ··· 10 ··· ··· 11 ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· 62 ABDO ET AL. Vol. 183 Table 5 (Continued) Name 0FGL R.A. Decl. l b θ95 100 MeV to 1 GeV 1 GeV to 100 GeV √ √ √ TS F23 ∆F23 TS23 F35 ∆F35 TS35 Var. J1310.6+3220 197.656 32.339 85.458 83.331 0.103 27.3 15.5 1.1 22.8 0.93 0.11 19.2 J1311.9−3419 197.998 −34.318 307.754 28.361 0.204 12.5 14.5 0.0 9.9 0.66 0.11 10.9 J1325.4−4303 201.353 −43.062 309.501 19.376 0.304 12.4 21.9 2.4 14.9 0.32 0.09 5.7 J1326.6−5302 201.651 −53.047 308.277 9.460 0.191 13.2 16.9 0.0 J1328.8−5604 202.222 −56.079 308.174 6.411 0.142 10.6 16.0 4.3 J1331.7−0506 202.935 −5.112 321.247 56.320 0.278 14.3 10.4 1.2 10.3 8.6 14.6 0.69 0.12 0.66 0.13 0.34 0.08 9.5 8.8 7.4 J1333.3+5058 203.331 50.973 107.300 64.865 0.219 12.4 8.8 J1355.0−1044 208.764 −10.735 327.221 49.113 0.163 11.5 8.4 1.3 1.3 11.5 9.7 0.28 0.07 0.30 0.07 7.9 8.0 J1413.1−6203 213.292 −62.063 312.346 −0.695 0.096 16.8 47.1 21.5 12.8 2.59 0.30 13.0 J1418.8−6058 214.718 −60.979 313.338 0.113 0.074 25.8 46.0 22.9 J1427.1+2347 216.794 23.785 29.472 68.166 0.073 24.1 4.2 1.0 18.3 10.1 5.42 0.38 0.92 0.11 22.1 21.4 J1430.5−5918 217.634 −59.301 315.288 1.173 0.119 10.4 29.7 12.1 J1457.6−3538 224.407 −35.639 329.936 20.530 0.076 39.6 30.1 0.5 9.9 32.7 1.26 0.21 2.00 0.17 8.7 26.0 J1459.4−6056 224.874 −60.937 317.863 −1.833 0.119 15.9 25.0 11.8 J1504.4+1030 226.115 10.505 11.409 54.577 0.054 88.2 63.4 2.1 12.0 71.8 1.26 0.19 5.86 0.26 10.1 58.9 J1509.5−5848 227.390 −58.812 320.003 −0.596 0.121 12.8 45.9 0.0 J1511.2−0536 227.814 −5.613 354.099 42.948 0.252 10.8 8.3 1.7 7.9 8.9 2.11 0.26 0.34 0.08 11.4 7.1 J1512.7−0905 228.196 −9.093 351.282 40.153 0.087 45.0 48.8 2.3 44.0 1.84 0.16 23.8 J1514.3−4946 228.585 −49.769 325.254 6.807 0.120 11.0 14.9 0.0 J1517.9−2423 229.496 −24.395 340.724 27.521 0.101 12.3 4.8 0.6 7.2 7.5 0.71 0.12 0.39 0.08 9.8 10.5 J1522.2+3143 230.552 31.726 50.143 57.014 0.087 34.3 21.2 1.5 30.3 1.06 0.11 20.7 J1536.7−4947 234.197 −49.798 328.261 4.764 0.127 10.7 19.0 0.0 J1543.1+6130 235.784 61.504 95.383 45.370 0.160 10.5 3.0 1.4 5.7 5.4 0.74 0.13 0.28 0.06 10.3 9.5 J1553.4+1255 238.368 12.922 23.746 45.225 0.105 23.7 14.5 2.2 17.5 1.08 0.12 18.8 J1555.8+1110 238.951 11.181 21.911 43.941 0.054 31.5 8.7 2.0 13.1 1.46 0.13 29.3 J1604.0−4904 J1615.6−5049 J1622.4−4945 J1625.8−2527 0.0 9.6 9.0 1.3 7.0 13.9 13.6 10.9 0.96 2.46 3.25 0.64 0.18 0.34 0.36 0.12 10.9 10.1 11.9 9.0 J1625.9−2423 246.494 −24.393 353.005 16.995 0.257 10.1 10.1 0.9 J1634.9−4737 248.733 −47.632 336.839 −0.025 0.079 28.6 106.1 6.6 J1635.2+3809 248.821 38.158 61.118 42.333 0.116 27.3 16.7 1.3 9.8 28.2 23.4 0.63 0.14 4.50 0.39 0.92 0.11 6.4 16.6 18.5 J1641.4+3939 J1648.1−4606 J1653.4−0200 J1653.9+3946 16.0 13.6 6.8 6.8 0.49 1.62 0.52 0.60 0.08 0.29 0.10 0.09 11.5 7.4 9.5 17.9 J1709.7−4428 257.427 −44.475 343.106 −2.679 0.048 85.8 115.3 5.0 54.0 15.82 0.49 68.7 J1714.7−3827 258.685 −38.459 348.525 0.103 0.133 16.0 58.8 6.0 15.3 11.6 241.015 243.914 245.611 246.470 250.355 252.029 253.355 253.492 −49.080 −50.831 −49.765 −25.451 39.666 −46.112 −2.014 39.767 332.170 332.354 333.874 352.164 63.239 339.469 16.549 63.612 2.541 −0.010 −0.009 16.308 41.239 −0.712 24.962 38.841 0.078 0.233 0.179 0.150 0.159 0.176 0.158 0.068 11.5 15.6 16.0 11.4 17.7 14.4 10.9 19.0 24.3 46.3 52.4 19.1 12.2 52.4 5.6 2.7 1.4 10.0 1.6 0.8 2.25 0.30 γ -Ray Assoc. Class ID or Assoc. EGR J1256−0552 3C 279 1AGL J1256−0549 T ··· bzq CGRaBS J1310+3220 B2 1308+32 · · · 3EG J1314−3431 · · · ··· EGR J1314−3417 · · · 3EG J1324−4314 rdg BZU J1325−4301 NGC 5128, Cen A ··· ··· ··· ··· T ··· ··· ··· T ··· bzq CGRaBS J1332−0509 PKS 1329−049 ··· ··· bzq CLASS J1333+5057 T ··· bzq CRATES J1354−1041 PKS 1352−104 · · · EGRc J1414−6224 · · · ··· 1AGL J1412−6149 · · · 1AGL J1419−6055 PSR LAT PSR J1418−6058 ··· ··· bzb CRATES J1427+2347 PKS 1424+240 ··· ··· ··· ··· T 3EG J1500−3509 bzq CGRaBS J1457−3539 PKS 1454−354 ··· ··· PSR LAT PSR J1459−60 T ··· bzq CGRaBS J1504+1029 PKS 1502+106 · · · 1AGL J1506−5859 PSR PSR J1509−5850 ··· ··· bzq PKS 1508−05 BZQ J1510−0543 T 3EG J1512−0849 bzq PKS 1510−08 EGR J1512−0857 BZQ J1512−0905 1AGL J1511−0908 ··· ··· ··· ··· ··· ··· bzb CGRaBS J1517−2422 AP Lib T ··· bzq CGRaBS J1522+3144 TXS 1520+319 ··· ··· ··· ··· ··· ··· bzb CRATES J1542+6129 RXS J15429+6129 T ··· bzq CRATES J1553+1256 PKS 1551+130 ··· ··· bzb CGRaBS J1555+1111 PG 1553+11 ··· ··· ··· ··· ··· ··· † ··· · · · 1AGL J1624−4946 · · · ··· · · · 3EG J1626−2519 bzq CGRaBS J1625−2527 PKS 1622−253 · · · 3EG J1627−2419 bzu CRATES J1627−2426 ··· ··· ··· ··· T 3EG J1635+3813 bzq CGRaBS J1635+3808 4C +38.41 T EGR J1642+3940 bzq CLASS J1641+3935 ··· ··· † ··· · · · 3EG J1652−0223 · · · ··· ··· ··· bzb CGRaBS J1653+3945 Mrk 501 · · · 3EG J1710−4439 PSR PSR J1709−4429 EGR J1710−4435 1AGL J1709−4428 · · · 3EG J1714−3857 † ··· Ref. ··· ··· ··· ··· ··· ··· ··· ··· ··· 7 ··· ··· 12 6 13 ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· No. 1, 2009 FERMI/LAT BRIGHT SOURCE LIST 63 Table 5 (Continued) Name 0FGL R.A. Decl. l b θ95 100 MeV to 1 GeV 1 GeV to 100 GeV √ √ √ TS F23 ∆F23 TS23 F35 ∆F35 TS35 Var. γ -Ray Assoc. Class ID or Assoc. J1719.3+1746 259.830 17.768 39.553 28.080 0.076 23.3 5.2 1.2 9.8 1.14 0.12 21.6 T J1732.8−3135 J1741.4−3046 J1742.1−2054 J1746.0−2900 J1751.5+0935 4.7 5.4 3.1 5.2 1.9 14.6 8.0 15.5 30.1 17.4 3.89 2.00 1.31 7.92 1.00 0.33 0.31 0.17 0.47 0.12 17.3 9.0 11.2 24.6 16.7 · · · EGR J1732−3126 ··· ··· · · · 3EG J1741−2050 T 3EG J1746−2851 T ··· −0.312 0.107 20.1 74.2 0.0 17.8 4.51 0.38 18.0 1.3 9.2 0.37 0.07 9.6 J1802.6−3939 270.661 −39.660 352.453 −8.410 0.069 25.2 19.3 0.7 17.7 1.48 0.16 18.2 J1805.3−2138 271.329 −21.649 8.536 J1809.5−2331 272.399 −23.520 7.381 20.0 24.3 3.13 0.32 5.63 0.31 12.7 32.0 6.5 6.3 21.6 8.5 20.8 14.1 13.8 26.3 2.79 1.91 1.55 5.76 0.24 0.30 0.29 0.37 18.3 8.9 7.4 22.8 J1826.3−1451 276.595 −14.860 16.886 −1.323 0.109 18.1 67.6 0.0 J1830.3+0617 277.583 6.287 36.158 7.543 0.097 12.8 16.8 0.0 J1833.4−2106 278.370 −21.103 12.109 −5.693 0.117 21.0 39.0 3.4 12.8 4.9 20.3 2.47 0.27 0.75 0.11 1.02 0.14 13.4 13.0 11.7 J1834.4−0841 278.617 −8.693 23.269 −0.220 0.100 11.7 44.6 4.1 J1836.1−0727 279.041 −7.460 24.557 −0.025 0.217 10.4 25.8 3.6 J1836.2+5924 279.056 59.406 88.855 24.996 0.053 98.0 44.4 1.7 11.9 9.0 69.8 1.34 0.29 1.99 0.30 8.36 0.31 8.3 8.2 73.8 J1839.0−0549 279.775 −5.826 26.343 0.078 0.119 22.2 51.5 3.5 J1844.1−0335 281.036 −3.589 28.907 −0.017 0.148 10.1 25.4 12.2 J1847.8+3223 281.954 32.385 62.065 14.838 0.178 16.0 8.8 0.6 18.3 9.0 13.9 4.04 0.33 1.77 0.28 0.64 0.10 17.0 8.1 11.3 J1848.6−0138 282.157 −1.640 31.153 −0.123 0.160 11.9 27.2 3.6 J1849.4+6706 282.365 67.102 97.503 25.027 0.090 28.0 13.8 1.5 9.8 21.7 2.09 0.28 1.07 0.12 9.8 20.1 J1855.9+0126 J1900.0+0356 J1907.5+0602 J1911.0+0905 J1911.2−2011 3.4 3.8 3.3 4.9 0.8 32.9 13.9 16.9 11.6 17.3 6.93 1.17 3.74 2.02 0.87 0.39 0.25 0.29 0.26 0.12 25.6 5.9 20.0 13.8 12.3 J1923.0+1411 290.768 14.191 49.134 −0.397 0.080 23.0 40.9 4.9 J1923.3−2101 290.840 −21.031 17.205 −16.199 0.130 16.4 10.6 0.6 15.3 13.0 3.37 0.27 0.68 0.11 19.1 11.7 J1953.2+3249 J1954.4+2838 J1958.1+2848 J2000.2+6506 · · · 3EG J1800−2338 EGR J1800−2328 ··· ··· bzb CGRaBS J1800+7828 S5 1803+78 T 3EG J1800−3955 bzu BZU J1802−3940 EGR J1758−3923 ··· ··· ··· ··· · · · 3EG J1809−2328 PSR LAT PSR J1809−2332 EGR J1809−2322 T ··· PSR LAT PSR J1813−1246 ··· ··· † ··· · · · 1AGL J1824−1414 · · · ··· · · · 3EG J1826−1302 PSR LAT PSR J1826−1256 1AGL J1824−1414 · · · 1AGL J1824−1414 hxb LS 5039 T ··· ··· ··· · · · 3EG J1832−2110 bzq BZQ J1833−2103 MC 1830−211 ··· ··· † ··· ··· ··· ··· ··· · · · 3EG J1835+5918 PSR LAT PSR J1836+5925 EGR J1835+5919 1AGL J1836+5923 ··· ··· ··· ··· ··· ··· ··· ··· T ··· bzq CGRaBS J1848+3219 TXS 1846+322 ··· ··· ··· ··· T 1AGL J1846+6714 bzq CGRaBS J1849+6705 S4 1849+67 · · · 1AGL J1857+0136 † ··· ··· ··· ··· ··· · · · 1AGL J1908+0613 PSR LAT PSR J1907+06 ··· ··· † ··· T 3EG J1911−2000 bzq CGRaBS J1911−2006 EGR J1912−2000 PKS 1908−201 ··· ··· † ··· T ··· bzq CGRaBS J1923−2104 TXS 1920−211 ··· ··· PSR PSR J1952+3252 ··· ··· † ··· · · · 3EG J1958+2909 PSR LAT PSR J1958+2846 ··· ··· bzb CGRaBS J1959+6508 1ES 1959+650 ··· ··· ··· ··· ··· ··· bzb CGRaBS J2009−4849 PKS 2005−489 ··· ··· bzu CLASS J2017+0603 · · · 1AGL J2021+3652 PSR PSR J2021+3651 · · · 1AGL J2022+4032 PSR LAT PSR J2021+4044 T 3EG J2025−0744 bzq CRATES J2025−0735 1AGL J2026−0732 PKS 2022−07 ··· ··· ··· ··· · · · 3EG J2033+4118 PSR LAT PSR J2032+4127 1AGL J2032+4102 ··· ··· ··· ··· 263.212 265.355 265.540 266.506 267.893 −31.588 −30.773 −20.916 −29.005 9.591 356.287 357.959 6.437 359.988 34.867 J1801.6−2327 270.404 −23.459 6.540 0.920 −0.189 4.859 −0.111 17.614 0.087 0.197 0.140 0.068 0.095 18.6 11.5 16.4 36.0 23.1 47.6 24.7 22.7 117.3 16.8 J1802.2+7827 270.567 78.466 110.026 28.990 0.132 12.6 5.5 J1813.5−1248 J1814.3−1739 J1821.4−1444 J1825.9−1256 273.399 273.581 275.365 276.497 283.984 285.009 286.894 287.761 287.813 298.325 298.614 299.531 300.053 −12.801 −17.665 −14.740 −12.942 1.435 3.946 6.034 9.087 −20.186 17.238 2.384 0.092 24.5 13.048 −0.094 0.191 13.1 16.435 −0.216 0.173 11.0 18.539 −0.344 0.075 30.9 34.722 37.424 40.140 43.246 16.818 −0.347 −0.110 −0.821 −0.176 −13.266 0.078 0.290 0.076 0.068 0.129 18.0 11.3 10.5 15.8 88.9 46.6 47.1 32.6 17.9 2.8 0.0 0.0 1.3 11.8 9.4 4.3 7.1 1.59 1.26 1.29 0.53 0.17 0.18 0.18 0.09 16.0 10.3 10.6 14.2 J2001.0+4352 300.272 43.871 79.047 7.124 0.069 13.3 9.5 J2009.4−4850 302.363 −48.843 350.361 −32.607 0.132 10.9 2.5 0.0 0.6 6.9 4.3 0.78 0.12 0.48 0.09 12.3 10.5 J2017.2+0602 J2020.8+3649 J2021.5+4026 J2025.6−0736 0.0 3.9 5.1 2.0 4.9 0.56 0.09 35.7 6.28 0.32 54.8 10.60 0.40 45.4 2.20 0.17 12.1 34.1 49.7 30.5 J2027.5+3334 306.882 33.574 73.296 −2.849 0.118 11.2 15.2 0.0 J2032.2+4122 308.058 41.376 80.161 0.978 0.085 23.9 51.3 4.7 6.7 18.1 0.97 0.15 3.07 0.26 9.3 18.8 J2055.5+2540 313.895 25.673 70.660 −12.475 0.130 17.3 5.4 12.6 0.88 0.11 13.6 6.048 36.830 40.439 −7.611 68.750 2.733 0.089 65.300 0.375 0.110 65.850 −0.232 0.112 97.974 17.630 0.077 39.0 11.6 25.9 17.0 20.0 53.7 50.0 56.6 95.2 13.6 23.3 16.1 2.9 304.302 305.223 305.398 306.415 32.818 28.649 28.803 65.105 −0.165 0.186 19.2 73.7 7.5 −1.938 0.061 39.8 41.1 4.0 48.596 −15.991 0.123 12.7 4.2 75.182 0.131 0.060 46.6 73.9 78.230 2.070 0.053 69.6 123.6 36.883 −24.389 0.077 50.6 40.8 0.8 ··· bzb CGRaBS J1719+1745 PKS 1717+177 PSR LAT PSR J1732−31 ··· ··· PSR LAT PSR J1741−2054 ··· ··· bzb CGRaBS J1751+0939 OT 081 † ··· Ref. ··· 7 ··· 7 ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· 7 ··· ··· ··· ··· ··· 7 ··· ··· ··· ··· ··· ··· ··· 7 ··· ··· ··· ··· ··· ··· 7 ··· ··· ··· ··· 14 7 ··· ··· 7 ··· 64 ABDO ET AL. Vol. 183 Table 5 (Continued) 100 MeV to 1 GeV Name 0FGL R.A. Decl. l b θ95 √ √ TS F23 ∆F23 TS23 1 GeV to 100 GeV √ F35 ∆F35 TS35 Var. J2056.1−4715 314.034 −47.251 352.586 −40.358 0.239 12.5 10.4 1.7 11.9 0.27 0.07 6.5 J2110.8+4608 317.702 46.137 88.261 −1.351 0.171 10.7 10.5 J2124.7−3358 321.186 −33.981 10.924 −45.441 0.143 14.3 3.5 J2139.4−4238 324.865 −42.642 358.237 −48.332 0.096 20.1 6.9 3.3 1.3 1.3 7.0 0.64 0.12 6.2 0.67 0.10 12.4 0.79 0.11 8.0 13.0 16.3 J2143.2+1741 325.807 17.688 72.016 −26.051 0.215 14.5 11.5 1.7 12.6 0.44 0.08 8.7 J2147.1+0931 326.777 9.519 1.6 19.1 0.55 0.09 12.1 6.9 1.5 8.6 0.38 0.08 7.3 J2158.8−3014 329.704 −30.237 17.711 −52.236 0.064 43.9 16.0 1.4 27.9 2.57 0.18 36.7 J2202.4+4217 330.622 42.299 92.569 −10.398 0.160 12.3 7.4 2.0 8.4 0.50 0.09 9.6 J2203.2+1731 330.815 17.532 75.715 −29.529 0.186 12.7 7.8 1.6 9.6 0.48 0.09 10.2 J2207.0−5347 331.765 −53.786 339.948 −49.832 0.223 12.4 10.8 1.8 12.0 0.26 0.07 6.6 J2214.8+3002 333.705 30.049 86.913 −21.658 0.152 11.9 4.2 J2229.0+6114 337.257 61.240 106.644 2.956 0.076 32.8 44.9 0.0 5.4 5.4 0.49 0.08 26.5 2.65 0.21 11.2 23.2 J2229.8−0829 337.452 −8.495 55.326 −51.701 0.185 16.8 11.7 0.4 16.1 0.30 0.07 7.5 J2232.4+1141 338.117 11.690 77.372 −38.592 0.183 15.2 10.8 1.3 13.2 0.35 0.07 8.2 J2241.7−5239 340.430 −52.651 337.395 −54.907 0.151 11.6 5.0 1.5 J2254.0+1609 343.502 16.151 86.125 −38.187 0.051 149.1 211.7 4.3 8.2 0.34 0.07 144.4 9.83 0.34 9.4 76.8 65.805 −32.236 0.137 19.9 16.0 J2157.5+3125 329.384 31.431 84.747 −18.258 0.343 10.0 J2302.9+4443 345.746 44.723 103.437 −14.004 0.155 13.6 J2325.3+3959 351.334 39.993 105.532 −19.952 0.118 11.4 J2327.3+0947 351.833 9.794 3.5 0.9 1.3 0.3 5.8 5.7 0.66 0.10 0.37 0.07 12.3 10.3 91.159 −47.821 0.218 17.1 15.5 1.6 17.1 0.34 0.07 8.4 0.0 1.3 5.5 0.55 0.09 15.9 0.36 0.08 12.7 8.1 J2339.8−0530 354.961 −5.512 81.487 −62.474 0.188 13.6 J2345.5−1559 356.389 −15.985 65.677 −71.092 0.239 15.5 4.9 9.9 γ -Ray Assoc. Class ID or Assoc. · · · 3EG J2055−4716 bzq CGRaBS J2056−4714 EGR J2057−4658 PKS 2052−47 ··· ··· ··· ··· ··· ··· PSR PSR J2124−3358 ··· ··· bzb CRATES J2139−4235 MH 2136−428 ··· ··· bzq CGRaBS J2143+1743 OX 169 T ··· bzq CGRaBS J2147+0929 PKS 2144+092 ··· ··· bzq CGRaBS J2157+3127 B2 2155+31 T 3EG J2158−3023 bzb CGRaBS J2158−3013 EGR J2200−3015 PKS 2155−304 · · · 3EG J2202+4217 bzb BZB J2202+4216 EGR J2204+4225 BL Lacertae T ··· bzq CGRaBS J2203+1725 PKS 2201+171 T ··· bzq CGRaBS J2207−5346 PKS 2204−54 ··· ··· ··· ··· · · · 3EG J2227+6122 PSR PSR J2229+6114 EGR J2227+6114 1AGL J2231+6109 ··· ··· bzq CGRaBS J2229−0832 PHL 5225 · · · 3EG J2232+1147 bzq BZQ J2232+1143 CTA 102 ··· ··· ··· ··· T 3EG J2254+1601 bzq CGRaBS J2253+1608 EGR J2253+1606 3C 454.3 1AGL J2254+1602 ··· ··· ··· ··· T ··· bzb CRATES J2325+3957 B3 2322+396 T ··· bzq CGRaBS J2327+0940 PKS 2325+093 ··· ··· ··· ··· T ··· bzq CGRaBS J2345−1555 PMN J2345−1555 Ref. ··· ··· 4 ··· ··· ··· ··· 15 ··· ··· ··· ··· 11 ··· ··· ··· 16 ··· ··· ··· ··· ··· Notes. Flux units 10−8 cm−2 s−1 . (†)—possible SNR or PWN association. See Table 2. A “0” in the ∆F23 column indicates that the entry in the F23 flux column is a 2σ upper limit. Reference. 1 (Abdo et al. 2008), “The Fermi Gamma-Ray Space Telescope Discovers the Pulsar in the Young Galactic Supernova Remnant CTA 1”; 2 (Abdo et al. 2009e), “Discovery of High-Energy Gamma-Ray Emission from the Globular Cluster 47 Tucanae with Fermi”; 3 (Abdo et al. 2009o), “Pulsed Gamma-Rays from the Millisecond Pulsar J0030+0451 with the Fermi Large Area Telescope”; 4 (Abdo et al. 2009d), “Discovery of a Population of Gamma-Ray Millisecond Pulsars with the Fermi Large Area Telescope”; 5 (Abdo et al. 2009j), “Fermi LAT Observations of LS I +61 303”; 6 (Abdo et al. 2009h), “Fermi Discovery of Gamma-Ray Emission from NGC1275”; 7 (Abdo et al. 2009p), “Sixteen Gamma-Ray Pulsars Discovered in Blind Frequency Searches Using the Fermi LAT”; 8 (Abdo et al. 2009b), “Fermi LAT Observations of the Vela Pulsar”; 9 (Abdo et al. 2009k), “Fermi/LAT Discovery of Gamma-ray Emission from a Relativistic Jet in the Narrow-line Quasar PMN J0948+0022”; 10 (Abdo et al. 2009f), “Discovery of Pulsed Gamma-Rays from the Young Radio Pulsar PSR J1028-5819 with the Fermi Large Area Telescope”; 11 (Abdo et al. 2009i), “Fermi LAT Detection of Pulsed Gamma-Rays from the Vela-like Pulsars PSR J1048-5832 and PSR J2229+6114”; 12 (Abdo et al. 2009l), “Fermi/LAT Discovery of Gamma-Ray Emission from the Flat-Spectrum Radio Quasar PKS 1454–354”; 13 (Abdo et al. 2009m), “PKS 1502+106: A New and Distant Gamma-Ray Blazar in Outburst Discovered by the Fermi Large Area Telescope”; 14 (Abdo et al. 2009n), “Pulsed Gamma-Rays from PSR J2021+3651 with the Fermi Large Area Telescope”; 15 (Aharonian et al. 2009), “Resolving the Blazar High-Energy Spectrum of PKS 2155-304 with HESS and Fermi”; 16 (Abdo et al. 2009g), “Early Fermi Gamma-Ray Space Telescope Observations of the Blazar 3C 454.3.” 4. Considering the highest confidence sources, in its lifetime EGRET found 31 sources (in either the 3EG or EGR catalogs or both) with confidence level of >10σ . The 0FGL list shows the dramatic improvement in sensitivity of the LAT. 5. Five of the EGRET sources seen at 10σ significance (all associated with flaring blazars: NRAO 190, NRAO 530, 1611+343, 1406−076, and 1622−297) do not appear in the LAT bright source list. 6. Twenty-eight of the EGRET sources that have counterparts in the 0FGL list were previously listed as unidentified. Half of these, 14, have now been firmly identified in this early LAT analysis. Thirteen are pulsars; 1 is a HMXB. No. 1, 2009 FERMI/LAT BRIGHT SOURCE LIST Table 6 LAT Bright Source List Source Associations (Firm Identifications) Class Number Radio/X-ray pulsar (PSR) LAT gamma-ray pulsar (LAT PSR) HMXB BL LAC (bzb) FSRQ (bzq) Other blazar (Uncertain type, bzu) Radio galaxy (rdg) Globular Cluster (glb, see the text) LMC (see the text) † Special cases (see Table 2) Unassociated 15 (15) 15 (15) 2 (1) 46 (0) 64 (0) 9 (0) 2 (0) 1 (0) 1 (0) 13 (0) 37 (0) 7. Of the 40 sources in the first AGILE catalog (which is contemporaneous but does not overlap in time with the 0FGL data), 32 are also found in the 0FGL list and seven more, while not formally overlapping, are “near misses” to 0FGL sources. The one exception is AGL J1238+04, a source associated with a FSRQ. A LAT source consistent in position with this one is found at a lower significance during the first three months of operation, but the source has since flared (Tramacere & Rea 2009). 6.3. Some Results from the Association Analysis Table 6 summarizes the census of associations in the bright source list. The numbers of these associations that are considered firm identifications are shown in parentheses. 1. The AGN class (121 members) is the largest source type associated in the LAT data. Details of the analysis, together with the implications for AGN studies, are given by Abdo et al. (2009c). Two of the AGNs found in this analysis are associated with radio galaxies; the rest are categorized as blazars. Note that five of the 0FGL AGNs are not included in the Abdo et al. (2009c) analysis because they are found within 10◦ of the Galactic plane. 2. Pulsars, including young radio pulsars, millisecond radio pulsars, and radio-quiet pulsars, form another well defined class (30 members) in the LAT bright source list. 3. Among the 0FGL sources, no associations were found with LMXB, starburst galaxies, prominent clusters of galaxies, or Seyfert galaxies. 4. Two associations were found with HMXB sources, both of which are also seen at TeV energies: LSI +61 303 (Albert et al. 2008) and LS 5039 (Aharonian et al. 2006b). The association of 0FGL J0240.3+6113 with LSI +61 303 is considered a firm identification based on the orbital periodicity seen in the LAT emission (Abdo et al. 2009j). Analysis of LS 5039 is in progress. 5. Globular cluster NGC 104 = 47 Tuc is associated with LAT source 0FGL J0025.1−7202; it should be emphasized that this globular cluster contains at least 23 millisecond radio pulsars and presumably contains many more as-yet undetected neutron stars. 6. 0FGL J0538.4−6856 is seen in the direction of the Large Magellanic Cloud. The LMC X-ray pulsar associated with this source by the automated software (PSR J0537−6910) is one possibility. The source is also consistent with the direction of the 30 Doradus star-forming region. Work on this part of the sky is still in progress. 65 7. 0FGL J0617.4+2234 lies within the projected direction of the shell of SNR IC443. A TeV source has also been seen close to the position of the LAT source. Detailed analysis of the LAT source is in progress. 8. 0FGL J0910.2−5044, although visible in the summed map, was seen primarily as a Galactic transient in 2008 October (Cheung et al. 2008). 9. 0FGL J1746.0−2900 lies close to the Galactic Center. Modeling the diffuse emission in this general region is challenging. We consider any conclusions about the association of this source with the Galactic Center or other candidate γ ray emitters in this region to be premature. The variability flag for this source is true, but the source is not extremely variable. This source barely met the criterion for being called variable. Work on this region is in progress. 10. Thirty-seven of the 0FGL sources have no obvious counterparts at other wavelengths. 6.4. TeV Comparisons Associations with TeV sources are based in this work only on positional correlation. Physical modeling or correlated variability would be needed in order to draw any conclusions from these associations. This is not an exhaustive list. We have omitted associations with blazars, well known objects such as the Crab Nebula, and sources discussed previously in the text. 1. 0FGL J1024.0−5754 is spatially consistent with HESS J1023−575, itself not yet firmly identified, but noted for its possible connection to the young stellar cluster Westerlund 2 in the star-forming region RCW49 (Aharonian et al. 2007). 2. 0FGL J1418.8−6058 is spatially coincident with HESS J1418−609 (Aharonian et al. 2006a), which may be the PWN powered by the LAT-discovered LAT PSR J1418−60. 3. 0FGL J1615.6−5049 is spatially coincident with HESS J1616-508, which has been suspected of being the PWN of PSR J1617−5055 (Landi et al. 2007; Kargaltsev et al. 2009). 4. 0FGL J1741.4−3046 is spatially consistent with the unidentified HESS J1741−302. (Tibolla et al. 2008). See the note in the previous section about LAT analysis in the Galactic Center region. 5. 0FGL J1805.3−2138 is spatially coincident with HESS J1804−216 (Aharonian et al. 2005). Formally still unidentified, HESS J1804−216 has been noted for possible counterparts in SNR G8.7−0.1, W30, or PSR J1803−2137. At this stage, we are not able to make a firm identification of the LAT source with any of the counterpart hypotheses. 6. 0FGL J1814.3−1739 is spatially coincident with HESS J1813−178 (Aharonian et al. 2005), which has been classified as a composite SNR, characterized by a shell-type SNR with central PWN candidate, not to be distinguishable given the angular resolution of present VHE observatories. At this stage, we are not able to settle either on a SNR or on a PSR/ PWN scenario for connecting HESS J1813−178 with the LAT source, leaving this study to a follow-up investigation. 7. 0FGL J1834.4−0841 is spatially coincident with HESS J1834−087 (Aharonian et al. 2005). Formally still unidentified, HESS J1834−087 was proposed to be explained in emission scenarios involving SNR W41, hadronic interactions with a giant molecular cloud, and/ or PSR J1833−0827. See Table 2. PSR J1833−0827 is not consistent in position with the LAT source. At this stage, we 66 ABDO ET AL. are not able to draw conclusions on a possible connection of the LAT source to the presented counterpart hypothesis. 8. 0FGL J1923.0+1411 is spatially coincident with HESS J1923+141, which is also spatially consistent with SNR G49.2−0.7 (W51). See Section 4.2.3 and Table 2 for a discussion of LAT source associations with SNRs and PWNe. 9. 0FGL J2032.2+4122, conclusively identified as a PSR, is spatially coincident with TeV 2032+4130 seen by HEGRA (Aharonian et al. 2002) and Milagro source MGRO J2031+41 (Abdo et al. 2007). Formally still unidentified, TeV J2032+4130 was noted for being close to the direction of the massive stellar cluster association Cygnus OB2. MGRO J2031+41, also unidentified, was reported as an extended and possibly confused source that could only be explained in part by the emission from TeV 2032+4130. We leave the possible association of LAT PSR J2032+41 with TeV J2032+4130 or MGRO 2031+41 to a detailed subsequent study. 10. Finally, it is noteworthy that LAT pulsars are found also in the error circles of four Milagro detected or candidate sources (Abdo et al. 2007): (a) 0FGL J2020.8+3649 (MGRO 2019+37), (b) 0FGL J1907.5+6002 (MGRO 1908+06), (c) 0FGL J0634.0+1745 (C3—candidate, Geminga), (d) 0FGL J2229.0+6114 (C4—candidate). Detailed discussion of individual sources is beyond the scope of this paper. By noting these positional coincidences, we call attention to areas of work still in progress on the Fermi/LAT data. The 0FGL list is a starting point for additional research in many areas. The Fermi/LAT Collaboration acknowledges generous ongoing support from a number of agencies and institutes that have supported both the development and the operation of the LAT as well as scientific data analysis. These include the National Aeronautics and Space Administration and the Department of Energy in the United States, the Commissariat à l’Energie Atomique and the Centre National de la Recherche Scientifique/Institut National de Physique Nucléaire et de Physique des Particules in France, the Agenzia Spaziale Italiana and the Istituto Nazionale di Fisica Nucleare in Italy, the Ministry of Education, Culture, Sports, Science and Technology (MEXT), High Energy Accelerator Research Organization (KEK), and Japan Aerospace Exploration Agency (JAXA) in Japan, and the K. A. Wallenberg Foundation, the Swedish Research Council, and the Swedish National Space Board in Sweden. Additional support for science analysis during the operations phase from the following agencies is also gratefully acknowledged: the Istituto Nazionale di Astrofisica in Italy and the K. A. Wallenberg Foundation in Sweden for providing a grant in support of a Royal Swedish Academy of Sciences Research fellowship for J.C. 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