The First Fermi Large Area Telescope Catalog of Gamma-ray Pulsars
aa r X i v : . [ a s t r o - ph . H E ] D ec Ap J Suppl 187, 460 (2010)
The First Fermi Large Area Telescope Catalog of Gamma-ray Pulsars
A. A. Abdo , , M. Ackermann , M. Ajello , W. B. Atwood , M. Axelsson , , L. Baldini ,J. Ballet , G. Barbiellini , , M. G. Baring , D. Bastieri , , B. M. Baughman , K. Bechtol ,R. Bellazzini , B. Berenji , R. D. Blandford , E. D. Bloom , E. Bonamente , , A. W. Borgland ,J. Bregeon , A. Brez , M. Brigida , , P. Bruel , T. H. Burnett , S. Buson ,G. A. Caliandro , , , R. A. Cameron , F. Camilo , P. A. Caraveo , J. M. Casandjian ,C. Cecchi , , ¨O. C¸ elik , , , E. Charles , A. Chekhtman , , C. C. Cheung , J. Chiang ,S. Ciprini , , R. Claus , I. Cognard , J. Cohen-Tanugi , L. R. Cominsky , J. Conrad , , ,R. Corbet , , S. Cutini , P. R. den Hartog , C. D. Dermer , A. de Angelis , A. de Luca , ,F. de Palma , , S. W. Digel , M. Dormody , E. do Couto e Silva , P. S. Drell , R. Dubois ,D. Dumora , , C. Espinoza , C. Farnier , C. Favuzzi , , S. J. Fegan , E. C. Ferrara , ,W. B. Focke , P. Fortin , M. Frailis , P. C. C. Freire , Y. Fukazawa , S. Funk , P. Fusco , ,F. Gargano , D. Gasparrini , N. Gehrels , , S. Germani , , G. Giavitto , B. Giebels ,N. Giglietto , , P. Giommi , F. Giordano , , T. Glanzman , G. Godfrey , E. V. Gotthelf ,I. A. Grenier , M.-H. Grondin , , J. E. Grove , L. Guillemot , , S. Guiriec , C. Gwon ,Y. Hanabata , A. K. Harding , M. Hayashida , E. Hays , R. E. Hughes , M. S. Jackson , , ,G. J´ohannesson , A. S. Johnson , R. P. Johnson , T. J. Johnson , , W. N. Johnson ,S. Johnston , T. Kamae , G. Kanbach , V. M. Kaspi , H. Katagiri , J. Kataoka , ,N. Kawai , , M. Kerr , J. Kn¨odlseder , M. L. Kocian , M. Kramer , , M. Kuss , J. Lande ,L. Latronico , M. Lemoine-Goumard , , M. Livingstone , F. Longo , , F. Loparco , ,B. Lott , , M. N. Lovellette , P. Lubrano , , A. G. Lyne , G. M. Madejski , A. Makeev , ,R. N. Manchester , M. Marelli , M. N. Mazziotta , W. McConville , , J. E. McEnery ,S. McGlynn , , C. Meurer , , P. F. Michelson , T. Mineo , W. Mitthumsiri , T. Mizuno ,A. A. Moiseev , , C. Monte , , M. E. Monzani , A. Morselli , I. V. Moskalenko , S. Murgia ,T. Nakamori , P. L. Nolan , J. P. Norris , A. Noutsos , E. Nuss , T. Ohsugi , N. Omodei ,E. Orlando , J. F. Ormes , M. Ozaki , D. Paneque , J. H. Panetta , D. Parent , , ,V. Pelassa , M. Pepe , , M. Pesce-Rollins , F. Piron , T. A. Porter , S. Rain`o , ,R. Rando , , S. M. Ransom , P. S. Ray , M. Razzano , N. Rea , , A. Reimer , ,O. Reimer , , T. Reposeur , , S. Ritz , A. Y. Rodriguez , R. W. Romani , , M. Roth ,F. Ryde , , H. F.-W. Sadrozinski , D. Sanchez , A. Sander , P. M. Saz Parkinson ,J. D. Scargle , T. L. Schalk , A. Sellerholm , , C. Sgr`o , E. J. Siskind , D. A. Smith , ,P. D. Smith , G. Spandre , P. Spinelli , , B. W. Stappers , J.-L. Starck , E. Striani , ,M. S. Strickman , A. W. Strong , D. J. Suson , H. Tajima , H. Takahashi , T. Takahashi ,T. Tanaka , J. B. Thayer , J. G. Thayer , G. Theureau , D. J. Thompson , S. E. Thorsett ,L. Tibaldo , , , O. Tibolla , D. F. Torres , , G. Tosti , , A. Tramacere , , Y. Uchiyama , ,T. L. Usher , A. Van Etten , V. Vasileiou , , , C. Venter , , N. Vilchez , V. Vitale , ,A. P. Waite , P. Wang , N. Wang , K. Watters , P. Weltevrede , B. L. Winer , K. S. Wood ,T. Ylinen , , , M. Ziegler Corresponding authors: G. A. Caliandro, [email protected]; E. C. Ferrara, [email protected]; D. Parent, [email protected]; R. W. Romani, [email protected]. Space Science Division, Naval Research Laboratory, Washington, DC 20375, USA National Research Council Research Associate, National Academy of Sciences, Washington, DC 20001, USA W. W. Hansen Experimental Physics Laboratory, Kavli Institute for Particle Astrophysics and Cosmology, De-partment of Physics and SLAC National Accelerator Laboratory, Stanford University, Stanford, CA 94305, USA Santa Cruz Institute for Particle Physics, Department of Physics and Department of Astronomy and Astrophysics,University of California at Santa Cruz, Santa Cruz, CA 95064, USA Department of Astronomy, Stockholm University, SE-106 91 Stockholm, Sweden The Oskar Klein Centre for Cosmoparticle Physics, AlbaNova, SE-106 91 Stockholm, Sweden Istituto Nazionale di Fisica Nucleare, Sezione di Pisa, I-56127 Pisa, Italy Laboratoire AIM, CEA-IRFU/CNRS/Universit´e Paris Diderot, Service d’Astrophysique, CEA Saclay, 91191 Gifsur Yvette, France Istituto Nazionale di Fisica Nucleare, Sezione di Trieste, I-34127 Trieste, Italy Dipartimento di Fisica, Universit`a di Trieste, I-34127 Trieste, Italy Rice University, Department of Physics and Astronomy, MS-108, P. O. Box 1892, Houston, TX 77251, USA Istituto Nazionale di Fisica Nucleare, Sezione di Padova, I-35131 Padova, Italy Dipartimento di Fisica “G. Galilei”, Universit`a di Padova, I-35131 Padova, Italy Department of Physics, Center for Cosmology and Astro-Particle Physics, The Ohio State University, Columbus,OH 43210, USA Istituto Nazionale di Fisica Nucleare, Sezione di Perugia, I-06123 Perugia, Italy Dipartimento di Fisica, Universit`a degli Studi di Perugia, I-06123 Perugia, Italy Dipartimento di Fisica “M. Merlin” dell’Universit`a e del Politecnico di Bari, I-70126 Bari, Italy Istituto Nazionale di Fisica Nucleare, Sezione di Bari, 70126 Bari, Italy Laboratoire Leprince-Ringuet, ´Ecole polytechnique, CNRS/IN2P3, Palaiseau, France Department of Physics, University of Washington, Seattle, WA 98195-1560, USA Columbia Astrophysics Laboratory, Columbia University, New York, NY 10027, USA INAF-Istituto di Astrofisica Spaziale e Fisica Cosmica, I-20133 Milano, Italy NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA Center for Research and Exploration in Space Science and Technology (CRESST), NASA Goddard Space FlightCenter, Greenbelt, MD 20771, USA University of Maryland, Baltimore County, Baltimore, MD 21250, USA George Mason University, Fairfax, VA 22030, USA Laboratoire de Physique et Chemie de l’Environnement, LPCE UMR 6115 CNRS, F-45071 Orl´eans Cedex 02, and Station de radioastronomie de Nan¸cay, Observatoire de Paris, CNRS/INSU, F-18330 Nan¸cay, France Laboratoire de Physique Th´eorique et Astroparticules, Universit´e Montpellier 2, CNRS/IN2P3, Montpellier,France Department of Physics and Astronomy, Sonoma State University, Rohnert Park, CA 94928-3609, USA Department of Physics, Stockholm University, AlbaNova, SE-106 91 Stockholm, Sweden Royal Swedish Academy of Sciences Research Fellow, funded by a grant from the K. A. Wallenberg Foundation Agenzia Spaziale Italiana (ASI) Science Data Center, I-00044 Frascati (Roma), Italy Dipartimento di Fisica, Universit`a di Udine and Istituto Nazionale di Fisica Nucleare, Sezione di Trieste, GruppoCollegato di Udine, I-33100 Udine, Italy Istituto Universitario di Studi Superiori (IUSS), I-27100 Pavia, Italy Universit´e de Bordeaux, Centre d’´Etudes Nucl´eaires Bordeaux Gradignan, UMR 5797, Gradignan, 33175, France CNRS/IN2P3, Centre d’´Etudes Nucl´eaires Bordeaux Gradignan, UMR 5797, Gradignan, 33175, France Jodrell Bank Centre for Astrophysics, School of Physics and Astronomy, The University of Manchester, M139PL, UK Arecibo Observatory, Arecibo, Puerto Rico 00612, USA Department of Physical Sciences, Hiroshima University, Higashi-Hiroshima, Hiroshima 739-8526, Japan University of Maryland, College Park, MD 20742, USA Istituto Nazionale di Fisica Nucleare, Sezione di Trieste, and Universit`a di Trieste, I-34127 Trieste, Italy University of Alabama in Huntsville, Huntsville, AL 35899, USA Department of Physics, Royal Institute of Technology (KTH), AlbaNova, SE-106 91 Stockholm, Sweden Australia Telescope National Facility, CSIRO, Epping NSW 1710, Australia Max-Planck Institut f¨ur extraterrestrische Physik, 85748 Garching, Germany Department of Physics, McGill University, Montreal, PQ, Canada H3A 2T8 Department of Physics, Tokyo Institute of Technology, Meguro City, Tokyo 152-8551, Japan Waseda University, 1-104 Totsukamachi, Shinjuku-ku, Tokyo, 169-8050, Japan Cosmic Radiation Laboratory, Institute of Physical and Chemical Research (RIKEN), Wako, Saitama 351-0198,Japan Centre d’´Etude Spatiale des Rayonnements, CNRS/UPS, BP 44346, F-30128 Toulouse Cedex 4, France Max-Planck-Institut f¨ur Radioastronomie, Auf dem H¨ugel 69, 53121 Bonn, Germany IASF Palermo, 90146 Palermo, Italy Istituto Nazionale di Fisica Nucleare, Sezione di Roma “Tor Vergata”, I-00133 Roma, Italy Department of Physics and Astronomy, University of Denver, Denver, CO 80208, USA Institute of Space and Astronautical Science, JAXA, 3-1-1 Yoshinodai, Sagamihara, Kanagawa 229-8510, Japan National Radio Astronomy Observatory (NRAO), Charlottesville, VA 22903, USA
ABSTRACT
The dramatic increase in the number of known gamma-ray pulsars since the launchof the
Fermi Gamma-ray Space Telescope (formerly GLAST) offers the first opportu-nity to study a sizable population of these high-energy objects. This catalog summa-rizes 46 high-confidence pulsed detections using the first six months of data taken bythe Large Area Telescope (LAT),
Fermi ’s main instrument. Sixteen previously un-known pulsars were discovered by searching for pulsed signals at the positions of brightgamma-ray sources seen with the LAT, or at the positions of objects suspected to beneutron stars based on observations at other wavelengths. The dimmest observed fluxamong these gamma-ray-selected pulsars is 6 . × − ph cm − s − (for E >
100 MeV).Pulsed gamma-ray emission was discovered from twenty-four known pulsars by usingephemerides (timing solutions) derived from monitoring radio pulsars. Eight of thesenew gamma-ray pulsars are millisecond pulsars. The dimmest observed flux among theradio-selected pulsars is 1 . × − ph cm − s − (for E >
100 MeV). The remaining sixgamma-ray pulsars were known since the
Compton Gamma Ray Observatory mission,or before. The limiting flux for pulse detection is non-uniform over the sky owing todifferent background levels, especially near the Galactic plane. The pulsed energy spec-tra can be described by a power law with an exponential cutoff, with cutoff energies inthe range ∼ − E ) of these neutron stars spans5 decades, from ∼ × erg s − to 5 × erg s − , and the apparent efficienciesfor conversion to gamma-ray emission range from ∼ .
1% to ∼ unity, although dis-tance uncertainties complicate efficiency estimates. The pulse shapes show substantial Institut de Ciencies de l’Espai (IEEC-CSIC), Campus UAB, 08193 Barcelona, Spain Sterrenkundig Institut “Anton Pannekoek”, 1098 SJ Amsterdam, Netherlands Institut f¨ur Astro- und Teilchenphysik and Institut f¨ur Theoretische Physik, Leopold-Franzens-Universit¨at Inns-bruck, A-6020 Innsbruck, Austria Space Sciences Division, NASA Ames Research Center, Moffett Field, CA 94035-1000, USA NYCB Real-Time Computing Inc., Lattingtown, NY 11560-1025, USA Dipartimento di Fisica, Universit`a di Roma “Tor Vergata”, I-00133 Roma, Italy Department of Chemistry and Physics, Purdue University Calumet, Hammond, IN 46323-2094, USA Max-Planck-Institut f¨ur Kernphysik, D-69029 Heidelberg, Germany Instituci´o Catalana de Recerca i Estudis Avan¸cats, Barcelona, Spain Consorzio Interuniversitario per la Fisica Spaziale (CIFS), I-10133 Torino, Italy North-West University, Potchefstroom Campus, Potchefstroom 2520, South Africa National Astronomical Observatories-CAS, ¨Ur¨umqi 830011, China School of Pure and Applied Natural Sciences, University of Kalmar, SE-391 82 Kalmar, Sweden & . Subject headings: catalogs – gamma rays: observations – pulsars: general – stars: neu-tronVersion corrected for an Erratum sent to the Ap J, December 2010: In the original paper, anerror was made in accounting for the delay due to interstellar dispersion in the radio phasing ofPSR J1124 − δ ) to 0.11 ±
1. Introduction
Following the 1967 discovery of pulsars by Bell and Hewish (Hewish et al. 1968), Gold (1968)and Pacini (1968) identified these objects as rapidly rotating neutron stars whose observable emis-sion is powered by the slow-down of the rotation. With their strong electric, magnetic, and gravita-tional fields, pulsars offer an opportunity to study physics under extreme conditions. As endpointsof stellar evolution, these neutron stars, together with their associated supernova remnants (SNRs)and pulsar wind nebulae (PWNe), help probe the life cycles of stars.Over 1800 rotation-powered pulsars are now listed in the ATNF pulsar catalog (Manchester et al.2005) , as illustrated in Figure 1. The vast majority of these pulsars were discovered by radio tele-scopes. Small numbers of pulsars have also been seen in the optical band, with more in the X-raybands (see e.g. Becker 2009).In the high-energy gamma-ray domain ( ≥
30 MeV) the first indications for pulsar emissionwere obtained for the Crab pulsar by balloon-borne detectors (e.g. Browning et al. 1971), andconfirmed by the SAS-2 satellite (Kniffen et al. 1974), which also found gamma radiation from theVela pulsar (Thompson et al. 1975). The
COS-B satellite provided additional details about these Compton Gamma Ray Observatory ( CGRO ) expanded the number of gamma-ray pulsarsto at least seven, with six clearly seen by the
CGRO high-energy instrument, EGRET. This gamma-ray pulsar population allowed a search for trends, such as the increase in efficiency η = L γ / ˙ E withdecreasing values of the open field line voltage of the pulsar, first noted by Arons (1996), for gamma-ray luminosity L γ and spin-down luminosity ˙ E . A summary of gamma-ray pulsar results in the CGRO era is given by Thompson (2004).The third EGRET catalog (3EG; Hartman et al. 1999) included 271 sources of which ∼
170 re-mained unidentified. Determining the nature of these unidentified sources is one of the outstandingproblems in high-energy astrophysics. Many of them are at high Galactic latitude and are mostlikely active galactic nuclei or blazars. However, most of the sources at low Galactic latitudes ( | b | ≤ ◦ ) are associated with star-forming regions and hence may be pulsars, PWNe, SNRs, winds frommassive stars, or high-mass X-ray binaries (e.g. Kaaret & Cottam 1996; Yadigaroglu & Romani1997; Romero et al. 1999). A number of radio pulsars were subsequently discovered in EGRETerror boxes (e.g. Kramer et al. 2003), but gamma-ray pulsations in the archival EGRET data werenever clearly seen. Solving the puzzle of the unidentified sources will constrain pulsar emission mod-els: pulsar population synthesis studies, such as those by Cheng & Zhang (1998), Gonthier et al.(2002), and McLaughlin & Cordes (2000), indicate that the number of detectable pulsars in ei-ther EGRET or Fermi data, as well as the expected ratio of radio-loud and radio-quiet pulsars(Harding et al. 2007), strongly depends on the assumed emission model.The Large Area Telescope (LAT) on the
Fermi Gamma-ray Space Telescope has provided amajor increase in the known gamma-ray pulsar population, including pulsars discovered first ingamma-rays (Abdo et al. 2009c) and millisecond pulsars (MSPs) (Abdo et al. 2009b). The firstaim of this paper is to summarize the properties of the gamma-ray pulsars detected by
Fermi -LATduring its first six months of data taking. The second primary goal is to use this gamma-ray pulsarcatalog to address astrophysical questions such as:1. Are all the gamma-ray pulsars consistent with one type of emission model?2. How do the gamma-ray pulsars compare to the radio pulsars in terms of physical propertiessuch as age, magnetic field, spin-down luminosity, and other parameters?3. Are the trends suggested by the
CGRO pulsars confirmed by the LAT gamma-ray pulsars?4. Which of the LAT pulsars are associated with SNRs, PWNe, unidentified EGRET sources,or TeV sources?The structure of this paper is as follows: Section 2 describes the LAT and the pulsar dataanalysis procedures; Section 3 presents the catalog and derives some population statistics from our 7 –sample; Section 4 studies the LAT sensitivity for gamma-ray pulsar detection, while in Section 5the implications of our results are briefly discussed. Finally, our conclusions are summarized inSection 6.
2. Observations and Analysis
The
Fermi Gamma-ray Space Telescope was successfully launched on 2008 June 11, carry-ing two gamma-ray instruments: the LAT and the Gamma-ray Burst Monitor (GBM). The LAT,
Fermi ’s main instrument, is described in detail in Atwood et al. (2009), with early on-orbit per-formance reported in Abdo et al. (2009o). It is a pair-production telescope composed of a 4 × ∼ for E > θ ≃ . ◦ E − . of a point source, with E in GeV, leveling offto θ . . ◦ E >
10 GeV . Effective area, PSF, and energy resolution are tabulated into bins ofphoton energy and angle of incidence relative to the LAT axis. The tables are called “instrumentresponse functions, and are described in detail in Abdo et al. (2009o). This work uses the versioncalled
P6 v3 diffuse .Gamma-ray events recorded with the LAT have time stamps that are derived from a GPS-synchronized clock on board the
Fermi satellite. The accuracy of the time stamps relative to UTCis < µ s (Abdo et al. 2009o). The timing chain from the GPS-based satellite clock through thebarycentering and epoch folding software has been shown to be accurate to better than a few µ sfor binary orbits, and significantly better for isolated pulsars (Smith et al. 2008).The LAT field-of-view is about 2.4 sr. Nearly the entire first year in orbit has been dedi-cated to an all-sky survey, imaging the entire sky every two orbits, i.e. every 3 hours. Any givenpoint on the sky is observed roughly 1 / th of the time. The LAT’s large effective area and ex-cellent source localization coupled with improved cosmic-ray rejection led to the detection of 46gamma-ray pulsars in the first six months of LAT observations. These include the six gamma-ray pulsars clearly seen with EGRET (Thompson 2004), two young pulsars seen marginally withEGRET (Ramanamurthy et al. 1996; Kaspi et al. 2000), the MSP seen marginally with EGRET(Kuiper et al. 2000), PSR J2021+3651 discovered in gamma-rays by AGILE (Halpern et al. 2008),and some of the other pulsars also studied by
AGILE (Pellizzoni et al. 2009).During the LAT commissioning period, several configuration settings were tested that affectedthe LAT energy resolution and reconstruction. However, these changes had no effect on the LATtiming. Therefore, for the spectral analyses, the data were collected from the start of the
Fermi sky-survey observations (2008 August 4, shortly before the end of the commissioning period) until 8 –2009 February 1, while the light curve and periodicity test analysis starts from the first eventsrecorded by the LAT after launch (2008 June 25) and also extends through 2009 February 1.s
We have conducted two distinct pulsation searches of
Fermi
LAT data. One search uses theephemerides of known pulsars, obtained from radio and X-ray observations. The other methodsearches for periodicity in the arrival times of gamma-rays coming from the direction of neutronstar candidates (“blind period searches”). Both search strategies have advantages. The former issensitive to lower gamma-ray fluxes, and the comparison of phase-aligned pulse profiles at differentwavelengths is a powerful diagnostic of beam geometry. The blind period search allows for thediscovery of new pulsars with selection biases different from those of radio searches, such as, forexample, favoring pulsars with a broader range of inclinations between the rotation and magneticaxes.For each gamma-ray event (index i ), the topocentric gamma-ray arrival time recorded by theLAT is transferred to times at the solar-system barycenter t i by correcting for the position of Fermi in the solar-system frame. The rotation phase φ i ( t i ) of the neutron star is calculated from a timingmodel, such as a truncated Taylor series expansion, φ i ( t i ) = φ + j = N X j =0 f j × ( t i − T ) j +1 ( j + 1)! . (1)Here, T is the reference epoch of the pulsar ephemeris and φ is the pulsar phase at t = T . Thecoefficients f j are the rotation frequency derivatives of order j . The rotation period is P = 1 /f .Different timing models, described in detail in Edwards et al. (2006), can take into account variousphysical effects. Most germane to the present work is accurate φ i ( t i ) computations, even in thepresence of the rotational instabilities of the neutron star called “timing noise”. “Phase-folding” alight curve, or pulse profile, means filling a histogram with the fractional part of the φ i values. Anephemeris includes the pulsar coordinates necessary for barycentering, the f j and T values, andmay include parameters describing the pulsar proper motion, glitch epochs, and more. The radiodispersion measure (DM) is used to extrapolate the radio pulse arrival time to infinite frequency,and uncertainties in the DM translate to an uncertainty in the phase offset between the radio andgamma-ray peaks. The ATNF database (version 1.36, Manchester et al. 2005) lists 1826 pulsars, and more havebeen discovered and await publication (Figure 1). The LAT observes them continuously during 9 –the all-sky survey. Phase-folding the gamma rays coming from the positions of all of these pulsars(consistent with the energy-dependent LAT PSF) requires only modest computational resources.However, the best candidates for gamma-ray emission are the pulsars with high ˙ E , which oftenhave substantial timing noise. Ephemerides accurate enough to allow phase-folding into, at aminimum, 25-bin phase histograms can degrade within days to months. The challenge is to have contemporaneous ephemerides.We have obtained 762 contemporaneous pulsar ephemerides from radio observatories, and 5from X-ray telescopes, in two distinct groups. The first consists of 218 pulsars with high spin-downpower ( ˙ E > erg s − ) timed regularly as part of a campaign by a consortium of astronomersfor the Fermi mission, as described in Smith et al. (2008). With one exception (PSR J1124 − Fermi launch. Some results from the timing campaign can befound in Weltevrede et al. (2009).The second group is a sample of 544 pulsars from nearly the entire P − ˙ P plane (Figure 2)being timed for other purposes for which ephemerides were shared with the LAT team. Thesepulsars reduce possible bias of the LAT pulsar searches created by our focus on the high ˙ E samplerequiring frequent monitoring. Gamma-ray pulsations from six radio pulsars with ˙ E < erg s − were discovered in this manner, all of which are MSPs.Table 1 lists properties of the 46 detected gamma-ray pulsars. Five of the 46 pulsars, all MSPs,are in binary systems. The period first derivative ˙ P is corrected for the kinematic Shklovskii effect(Shklovskii 1970): ˙ P = ˙ P obs − µ P obs d/c , where µ is the pulsar proper motion, and d the distance.The correction is small except for a few MSPs (Abdo et al. 2009b). The characteristic age is τ c = P/ P and the spin-down luminosity is˙ E = 4 π I ˙ P P − , (2)taking the neutron star’s moment of inertia I to be 10 g cm . The magnetic field at the lightcylinder (radius R LC = cP/ π ) is B LC = I π ˙ Pc P ! / ≈ . × ( ˙ P P − ) / G . (3)Table 2 lists which observatories provided ephemerides for the gamma-ray pulsars. An “L”indicates that the pulsar was timed using LAT gamma rays, as described in the next Section.“P” is the Parkes radio telescope (Manchester 2008; Weltevrede et al. 2009). The majority ofthe Parkes observations were carried out at intervals of 4 – 6 weeks using the 20 cm Multibeamreceiver (Staveley-Smith et al. 1996) with a 256 MHz band centered at 1369 MHz. At 6 monthintervals observations were also made at frequencies near 0.7 and 3.1 GHz with bandwidths of 32and 1024 MHz respectively. The required frequency resolution to avoid dispersive smearing acrossthe band was provided by a digital spectrometer system. 10 –“N” is the Nan¸cay radio telescope (Theureau et al. 2005). Nan¸cay observations are carriedout every three weeks on average. The recent version of the BON backend is a GPU-based coher-ent dedispersor allowing the processing of a 128 MHz bandwidth over two complex polarizations(Cognard et al. 2009). The majority of the data are collected at 1398 MHz, while for MSPs inparticular, observations are duplicated at 2048 MHz to allow DM monitoring.“J” is the 76-m Lovell radio telescope at Jodrell Bank in the United Kingdom. Jodrell Bankobservations (Hobbs et al. 2004) were carried out at typical intervals of between 2 days and 10days in a 64-MHz band centered on 1404 MHz, using an analog filterbank to provide the frequencyresolution required to remove interstellar dispersive broadening. Occasionally, observations are alsocarried out in a band centered on 610 MHz to monitor the interstellar dispersion delay.“G” is the 100-m NRAO Green Bank Telescope (GBT). PSR J1833–1034 was observed monthlyat 0.8 GHz with a bandwidth of 48 MHz using the BCPM filter bank (Backer et al. 1997). The otherpulsars monitored at the GBT were observed every two weeks at a frequency of 2 GHz across a600 MHz band with the Spigot spectrometer (Kaplan et al. 2005). Individual integration timesranged between 5 minutes and 1 hour.Arecibo (”A”) observations of the very faint PSRs J1930+1852 and J2021+3651 are carriedout every two weeks, with the L-wide receiver (1100 to 1730 MHz). The back-ends used are theWideband Arecibo Pulsar Processor (WAPP) correlators (Dowd et al. 2000), each with a 100MHz-wide band. The antenna voltages are 3-level digitized and then auto-correlated with a totalof 512 lags, accumulated every 128 µ s, and written to disk as 16-bit sums. Processing includesFourier-transforms to obtain power spectra, which are then dedispersed and phase-folded. Theaverage pulse profiles are cross-correlated with a low-noise template profile to obtain topocentrictimes of arrival.“W” is the Westerbork Synthesis Radio Telescope, with which observations were made approx-imately monthly at central frequencies of 328, 382 and 1380 MHz with bandwidths of 10 MHz at thelower frequencies and 80 MHz at the higher frequencies. The PuMa pulsar backend (Voˆute et al.2002) was used to record all the observations. Folding and dedispersion were performed offline.The rms of the radio timing residuals for most of the solutions used in this paper is < . .
2% for five pulsars. This is adequate for the 50- or25-bin phase histograms used in this paper. The ephemerides used for this catalog will be availableon the
Fermi
Science Support Center data servers . http://fermi.gsfc.nasa.gov/ssc/
11 –
For all 16 of the pulsars found in the blind searches of the LAT data, we determined the timingephemerides used in this catalog directly from the LAT data as described below. In addition, fortwo other pulsars the LAT data provided the best available timing model. The first is the radio-quiet pulsar Geminga. Since Geminga is such a bright gamma-ray pulsar, it is best timed directlyusing gamma-ray observations. During the period between EGRET and
Fermi , occasional
XMM-Newton observations maintained the timing model (Jackson & Halpern 2005) but a substantiallyimproved ephemeris has now been derived from the LAT data (Abdo et al. 2009a). The second isPSR J1124 − L γ of a young pulsar can be several percent of ˙ E , the gamma-ray counting ratesare low. As an example, the LAT detects a gamma ray from the Crab pulsar approximately every500 rotations, when the Crab is well within the LAT’s field-of-view. Such sparse photon arrivalsmake periodicity searches difficult. Extensive searches for pulsations performed on EGRET data(Chandler et al. 2001; Ziegler et al. 2008) were just sensitive enough to detect the very brightpulsars Vela, Crab, and Geminga in a blind search, had they not already been known as pulsars.Blind periodicity searches of all other EGRET sources proved fruitless.By contrast, the improvements afforded by the LAT have enabled highly successful blindsearches for pulsars. In the first six months of operation, we discovered a total of 16 new pulsarsin direct pulsation searches of the LAT data (see e.g. Abdo et al. 2008, 2009c). A computation-ally efficient time-difference search technique made these searches possible (Atwood et al. 2006),enabling searches of hundreds of Fermi sources to be performed on a small computer cluster withonly a modest loss in sensitivity compared to fully coherent search techniques. Still, owing to thelarge number of frequency and frequency derivative trials required to search a broad parameterspace, the minimum gamma-ray flux needed for a statistically significant detection is considerablyhigher than the minimum flux needed for the phase-folding technique using a known ephemeris (asin Section 2.1, Eq. 1).We performed these blind searches on ∼
100 candidate sources identified before launch and onanother ∼
200 newly detected LAT sources. The parameter space covered by the blind searchesincluded frequencies from 0.5 Hz to 64 Hz (periods of 156.25 ms to 2 s), and a frequency derivativefrom zero to the spin-down of the young Crab pulsar ( f = − . × − ), which covers ∼
86% ofthe pulsars contained in the ATNF database (Abdo et al. 2009c). Of the 16 pulsars detected inthese searches, 13 are associated with previously known EGRET sources. The discoveries includeseveral long-suspected pulsars in SNRs and PWNe.These 16 pulsars are gamma-ray selected, as they were discovered by the LAT and thus thepopulation is subject to very different selection effects than the general radio pulsar population. 12 –However, this does not necessarily imply that they are radio quiet. For several cases, deep radiosearches have already been performed on known PWNe or X-ray point sources suspected of har-boring pulsars. In most cases, new radios searches are required to investigate whether there is aradio pulsar counterpart down to a meaningful luminosity limit. These searches are now beingundertaken and are yielding the first results (Camilo et al. 2009).For these 18 pulsars (16 new plus Geminga and PSR J1124 − < . ◦ or < ◦ (see further Section 2.1.3 and Table 2). For this pulsartiming analysis, we used diffuse class photons with energies above a cutoff (typically E >
300 MeV)selected to optimize the signal-to-noise ratio for that particular pulsar. We converted the photonarrival times to the geocenter using the gtbary science tool . This correction removes the effectsof the spacecraft motion about the Earth, resulting in times as would be observed by a virtualobservatory at the geocenter.Using an initial timing model for the pulsar, we then used Tempo2 (Hobbs et al. 2006) in itspredictive mode to generate polynomial coefficients describing the pulse phase as a function of timefor an observatory at the geocenter. Using these predicted phases, we produced folded pulse profilesover segments of the LAT observation. The length of the segments depends on the brightness of thepulsar but are typically 10–20 days. We then produced a pulse time of arrival (TOA) for each datasegment by Fourier domain cross-correlation with a template profile (Taylor 1992). The templateprofile for most of the pulsars is based on a multi-Gaussian fit to the observed LAT pulse profile.However, in the case of Geminga, which has very high signal-to-noise and a complex profile notwell described by a small number of Gaussians, we used a template profile that was the full missionlight curve itself.Finally, we used
Tempo2 to fit a timing model to each pulsar. For most of the pulsars, themodel includes pulsar celestial coordinates, frequency and frequency derivative. In several cases,the fit also required a frequency second derivative term to account for timing noise. In the caseof PSR J1124 − − − . For Geminga Pulsar timing positions are measured by fitting the sinusoidal delays of the pulse arrival times associated with theEarth moving along its orbit. For pulsars very close to the ecliptic plane the derivative of this delay with respect to
13 –and PSR J1124 − Fermi timing. The rms of the timing residuals are between 0 . .
9% of arotation period, with the highest being for PSR J1459 −
60. The
Tempo2 timing models used forthe catalog analysis will be made available online at the FSSC web site . The light curves of 46 gamma-ray pulsars detected by the LAT are appended to the end ofthis paper, in Figures A-1 to A-46. The gray light curve in the top panel includes all photons with
E > . E > . E > . . . . . N = 25 or 50 bins, with 25 bins usedwhen required to keep at least 50 counts per bin in the peak of the light curve or to prevent unduesmearing due to the accuracy of the timing model.Table 3 lists light curve parameters, taken from the >
100 MeV profiles (top panel of FiguresA-1–A-46). For some pulsars (e.g. PSR J1420 − Z (Buccheri et al. 1983) and H (de Jager et al. 1989) periodicity test valuesfor E > . ∼
800 pulsars searched. Detection of gamma-ray pulsationsare claimed when the significance of the periodicity test exceeds 5 σ (i.e. a chance probability of < × − ). We have used the Z -test with m = 2 harmonics ( Z ) which provides an analyticaldistribution function for the null hypothesis described by a χ distribution with 2 m degrees offreedom. The H -test uses Monte-Carlo simulations to calculate probabilities, limited to a minimumof 4 × − (equivalent to 5.37 σ ). Each method is sensitive to different pulse profile shapes. Fourpulsars in the catalog fall short of the 5 σ significance threshold in the six-month data set with theselection cuts applied here: the 3 MSPs J0218+4232, J0751+1807, and J1744 − . ◦
0, but 0 . ◦ ecliptic latitude is greatly reduced and thus such pulsars have one spatial dimension poorly constrained, as discussedin Ray et al. (2010).
14 –by using the energy spectrum for the phase-averaged source, described in Section 2.2, to maximize S /N over a grid of maximum radii and minimum energy thresholds (where S is the number ofcounts attributed to the point source, and N is the number of counts due to the diffuse backgroundand neighboring sources). We selected photons within a radius θ of the pulsar position, requiringa radius of at least 0 . ◦
35, but no larger than the reported maxROI.We estimated the background level represented by the dashed horizontal lines in Figures A-1to A-46 from an annulus between 1 ◦ < θ < ◦ surrounding the source. Nearby sources were removed,and we normalized to the same solid angle as the source ROI. The poor spatial resolution of theLAT at low energies can blur structured diffuse emission and bias this background estimate. Thelevels shown are intended only to guide the eye. Detailed analyses of off-pulse emission will bediscussed in future work. The pulsar spectra were fitted with an exponentially cutoff power-law model of the formd N d E = KE − ΓGeV exp (cid:18) − EE cutoff (cid:19) (4)in which the three parameters are the photon index at low energy Γ, the cutoff energy E cutoff , anda normalization factor K (in units of ph cm − s − MeV − ), in keeping with the observed spectralshape of bright pulsars (Abdo et al. 2009g). The energy at which the normalization factor K isdefined is arbitrary. We chose 1 GeV because it is, for most pulsars, close to the energy at whichthe relative uncertainty on the differential flux is minimal.We wish to extract the spectra down to 100 MeV in order to constrain the power-law part ofthe spectrum, and to measure the flux above 100 MeV directly. Because the spatial resolution of Fermi is not very good at low energies ( ∼ ◦ at 100 MeV), we need to account for all neighboringsources and the diffuse emission together with each pulsar. This was done using the frameworkused for the LAT Bright Source List (Abdo et al. 2009i). A 6-month source list was generated inthe same way as the 3-month source list described in Abdo et al. (2009i), but covering the extendedperiod of time used for the pulsar analysis. We have added the source Cyg X-3 (Abdo et al. 2009l),although it was not detected automatically as a separate source, because it is very close to PSRJ2032+4127, and impacts the spectral fit of the pulsar. Cyg X-3 was fit with a simple power-lawas were all other non-pulsar sources in the list.We used a Galactic diffuse model designated
54 77Xvarh7S calculated using GALPROP , anevolution of that used in Abdo et al. (2009i). A similar model, gll iem v02 , is publicly available .We kept events with E >
100 MeV belonging to the diffuse event class, which has the tightest http://galprop.stanford.edu/
15 –cosmic-ray background rejection (Atwood et al. 2009). To avoid contamination by gamma-raysproduced by cosmic-ray interactions in the Earth’s atmosphere, we select time intervals when theentire ROI, of radius 10 ◦ around the source, has a zenith angle < ◦ . We extracted events in acircle of radius 10 ◦ around each pulsar, and included all sources up to 17 ◦ into the model (sourcesoutside the extraction region can contribute at low energy). Sources further away than 3 ◦ fromthe pulsar were assigned fixed spectra, taken from the all-sky analysis. Spectral parameters for thepulsar and sources within 3 ◦ of it were left free for the analysis.The fit was performed by maximizing unbinned likelihood (direction and energy of each eventis considered) as described in Abdo et al. (2009i) and using the minuit fitting engine . The uncer-tainties on the parameters were estimated from the quadratic development of the log(likelihood)surface around the best fit. In addition to the index Γ and the cutoff energy E cutoff which areexplicit parameters of the fit, the important physical quantities are the photon flux F (in unitsof ph cm − s − ) and the energy flux G (in units of erg cm − s − ), F = Z
100 GeV100 MeV d N d E d E, and (5) G = Z
100 GeV100 MeV E d N d E d E. (6)These derived quantities are obtained from the primary fit parameters. Their statistical uncertain-ties are obtained using their derivatives with respect to the primary parameters and the covariancematrix obtained from the fitting process.For a number of pulsars, an exponentially cutoff power-law spectral model is not significantlybetter than a simple power-law. We identified these by computing T S cutoff = 2∆log(likelihood)(comparable to a χ distribution with one degree of freedom) between the models with and withoutthe cutoff. Pulsars with T S cutoff <
10 have poorly measured cutoff energies.
T S cutoff is reported inTable 4.The above analysis yields a fit to the overall spectrum, including both the pulsar and anyunpulsed emission, such as from a PWN. To do better we split the data into on-pulse and off-pulsesamples and modeled the off-pulse spectrum by a simple power-law. The off-pulse window used forthis background estimation is defined in the last column of Table 3.In a second step we re-fitted the on-pulse emission to the exponentially cutoff power-law asbefore, with the off-pulse emission (scaled to the on-pulse phase interval) added to the model andfixed to the off-pulse result. In many cases the off-pulse emission was not significant at the 5 σ oreven 3 σ level, but we kept the formal best fit anyway, in order to not bias the pulsed emissionupwards. The results summarized in Table 4 come from this on-pulse analysis.Using an off-pulse pure power law is not ideal for the Crab or any other PWN with synchrotron http://lcgapp.cern.ch/project/cls/work-packages/mathlibs/minuit/doc/doc.html
16 –and inverse Compton components within the
Fermi energy range. Judging from the Crab pulsar,using a simple power-law to model the off-pulse emission mainly affects the value of the cutoffenergy. The analysis specific to the Crab, with a model adapted to the pulsar synchrotron compo-nent low energies and to the high energy nebular component, yields E cutoff ∼ >
10 GeV). The photon and energy fluxes givenby the two analyses are within 10% of each other. Additional exceptions in Table 4 are for PSRsJ1836+5925 and J2021+4026. The off-pulse phase definition for these pulsars is unclear, so thespectral parameters reported in the Table are from the initial, phase-averaged spectral analysis.We have checked whether our imperfect knowledge of the Galactic diffuse emission may impactthe pulsar parameters by applying the same analysis with a different diffuse model, as was done inAbdo et al. (2009i). The phase-averaged emission is affected. Seven (relatively faint) pulsars seetheir flux move up or down by more than a factor of 1.5. On the other hand, the pulsed flux ismuch more robust, because the off-pulse component absorbs part of the background difference, andthe source-to-background ratio is better after on-pulse phase selection. Only two pulsars see theirpulsed flux move up or down by more than a factor of 1.2, and none shift by more than a factor of 1.4when changing the diffuse model. Overall, the systematic uncertainties due to the diffuse model onthe fluxes F and G , on the photon index Γ, and on E cutoff scale with the statistical uncertainty.Adding the statistical and systematic errors in quadrature amounts, to a good approximation, tomultiplying the statistical errors on F , G , and Γ by 1 .
2. The uncertainties listed in Table 4and plotted in the figures include this correction. The increase in the uncertainty on E cutoff due tothe diffuse model is <
5% and is neglected.Systematic uncertainties on the LAT effective area are of order 5% near 1 GeV, 10% below 0 . δ Γ =(+0 . , − . δE cutoff = (+20% , − δF = (+30% , − δG = (+20% , − . ◦
35 and a maximum of 5 ◦ . 17 –The observed pulsed spectrum was built by selecting the events in the on-pulse phase intervaland subtracting the events in the off-pulse interval, properly scaled for the phase ratio. Theinstrument response function, expressed as a smearing matrix, was evaluated using the LAT Geant4 -based Monte Carlo simulation package called
Gleam (Boinee et al. 2003), taking into account thepointing history of the source.The true pulsar energy spectra were then reconstructed from the observed ones using aniterative procedure based on Bayes’ theorem (Mazziotta 2009). Typically, convergence is reachedafter a few iterations. When the procedure has converged, both statistical and systematic errorson the observed energy distribution can be easily propagated to the unfolded spectra. The resultsobtained from the unfolding analysis were consistent within errors with the likelihood analysisresults.
3. The LAT Pulsar Sample
We describe here the astronomical context of the observed LAT pulsars, including our currentbest understanding of the source distances, the Galactic distribution and possible associations.We also note correlations among some observables which may help probe the origin of the pulsaremission.
Converting measured pulsar fluxes to radiated power requires reliable distance estimates. An-nual trigonometric parallax measurements are the most reliable, but are generally only availablefor a few relatively nearby pulsars.The most commonly used technique to obtain radio pulsar distances exploits the pulse de-lay as a function of wavelength by free electrons along the path to Earth. A distance can becomputed from the DM coupled to an electron density distribution model. We use the NE2001model (Cordes & Lazio 2002) unless noted otherwise. It assumes uniform electron densities in andbetween the Galactic spiral arms, with smooth transitions between zones, and spheres of greaterdensity for specific regions such as the Gum nebula, or surrounding Vela. Specific lines-of-sight cantraverse unmodeled regions of over- (or under-) density, as, for example, along the tangents of thespiral arms, causing significant discrepancies between the true pulsar distances and those inferredfrom the electron-column density.A third method, kinematic, associates the pulsar with objects whose distance can be measuredfrom the Doppler shift of absorption or emission lines in the neutral hydrogen (HI) spectrum, http://geant4.web.cern.ch/geant4/
18 –together with a rotation curve of the Galaxy. It breaks down where the velocity gradients arevery small or where the distance-velocity relation has double values. The associations are oftenuncertain, and these distance measurements can be controversial.In a small number of cases, the distance is evaluated either from X-ray measurements of theabsorbing column at low energies (below 1 keV), or from consideration of the X-ray flux assumingsome standard parameters for the neutron star.Table 5 presents the best known distances of 37 pulsars detected by
Fermi , the methods usedto obtain them, and the references. For distances obtained from the NE2001 model and the DM,the reference indicates the DM measurement. We assume a minimum DM distance uncertaintyof 30%. When distances from different methods disagree and no method is more convincing thanthe other, a range is given, and 30% uncertainties on the upper and lower values are used. Forthe remaining 9
Fermi -discovered pulsars no distance estimates have been established so far. Herefollow comments for some of the distance values reported in Table 5:
PSR J0205+6449 – The pulsar is in the PWN 3C 58. NE2001 gives 4.5 kpc for DM=141cm − pc in this direction (Camilo et al. 2002d). Using HI absorption and emission lines from thePWN yields from 2.6 kpc (Green & Gull 1982) to 3.2 kpc (Roberts et al. 1993). The lower V-band reddening (Fesen et al. 1988, 2008) compared to the Galactic-disk edge (Schlegel et al. 1998)suggests that the PWN is in the range 3–4 kpc. Table 5 quotes the distance range found byGreen & Gull (1982) and Roberts et al. (1993). PSR J0218+4232 – The DM measurements from Navarro et al. (1995) together with NE2001yield 2.7 ± PSR J0248+6021 – The DM of 376 cm − pc (Cognard et al. 2010) puts this pulsar beyondthe edge of the Galaxy for this line-of-sight. The line-of-sight, however, borders the giant HII regionW5 in the Perseus Arm. We bracket the pulsar distance as being between W5 (2 kpc) and theGalaxy edge (9 kpc). PSR J0631+1036 – The DM = 125 . − pc (Zepka et al. 1996) is large for a source in thedirection of the Galactic anticenter. The dark cloud LDN 1605, part of the active star-formingregion 3 Mon, is in the line-of-sight. The pulsar could be inside the cloud, at ∼ PSR J1124 − – It lies towards the Carina arm where NE2001 biases are acute. The DM dis-tance is 5.7 kpc (Camilo et al. 2002c). The kinematic distance of the associated SNR (G292.0+1.8)indicates a lower limit of 6.2 ± PSR J1418 − – This pulsar is likely associated with the PWN G313.3+0.1, near theKookaburra complex. A nearby HII region is at 13.4 kpc (Caswell & Haynes 1987) but could easily 19 –be in the background. Such a large distance implies an unreasonably large gamma-ray efficiency.Table 5 lists a crude estimate of the distance range with the lower limit (Yadigaroglu & Romani1997) taking the pulsar to be related to one of the near objects (Clust 3, Cl Lunga 2 or SNRG312.4 − PSR J1709 − – The NE2001 DM distance is 2.3 ± ± ± Chandra (Romani et al. 2005) and
XMM-Newton (McGowan et al. 2004) is compatible with a distance of 1.4–2.0 kpc. We assume the range1.4–3.6 kpc.
PSR J1747 − – The pulsar is associated with the PWN G359.23 − − ) suggests2.0 ± Chandra is between 4and 5 kpc, while the closer value of 2 kpc would imply that an otherwise unknown molecular cloudlies in front of the pulsar (Gaensler et al. 2004). A range of 2–5 kpc is used in our analysis.
PSR J2021+3651 – The DM distance of ∼
12 kpc implies a high gamma-ray conversion effi-ciency (Roberts et al. 2002; Abdo et al. 2009m). The open cluster Berkeley 87 near the line-of-sightcould be responsible for an electron column density higher than modeled by NE2001. The distancein Table 5 comes from a
Chandra
X-ray observation of the pulsar and its surrounding nebula(Van Etten et al. 2008). A similar range (1.3–4.1 kpc) was obtained for the X-ray flux detectedfrom the associated PWN.
PSR J2032+4127 – The DM value (115 pc cm − ) gives an NE2001 distance of 3.6 kpc. Ifthe pulsar belongs to the star cluster Cyg OB2, it would be located at approximately 1.6 kpc(Camilo et al. 2009). In this text we use a range of 1.6–3.6 kpc for this source. PSR J2229+6114 – The distance derived from the X-ray absorption is ∼ Figure 1 shows the pulsars projected on the sky. A Gaussian fit to the Galactic latitudedistribution for those with | b | < ◦ and having distance estimates yields a standard deviation of σ b = 3 . ± . d = 0 .
25 to 5 . h < d sin 3 . ◦ δ from Table 3 are summarized in Figure4. As we will discuss in Section 5, outer magnetosphere emission models predict correlations betweenthese parameters. Figure 5 shows B LC versus the characteristic age ( τ c ). The magnetic fields atthe light cylinder for the detected MSPs are comparable to those of the other gamma-ray pulsars,suggesting that the emission mechanism for the two families may be similar.Table 4 lists L γ and η , while Figure 6 plots L γ vs. ˙ E . The dashed line indicates L γ = ˙ E , whilethe dot-dashed line indicates L γ ∝ ˙ E / , where L γ ≡ πd f Ω G . (7)The flux correction factor f Ω (Watters et al. 2009) is model-dependent and depends on the magneticinclination and observer angles α and ζ . Both the outer gap and slot gap models predict f Ω ∼ f Ω = 1 / π ≈ .
08 (in e.g. Thompson et al. 1994), or f Ω = 0 . f Ω = 1 throughout the paper, whichpresumably induces an artificial spread in the quoted L γ values. However, it is the quadraticdistance dependence for L γ that dominates the uncertainty in L γ in nearly all cases.Gamma-rays dominate the total power L tot radiated by most known high-energy pulsars, thatis, L tot ≈ L γ . The Crab is a notable exception, with X-ray luminosity L X ∼ L γ . In Figure 6 weplot both L γ and L X + L γ . L X for E <
100 MeV is taken from Figure 9 of Kuiper et al. (2001).
Table 6 provides some alternate names and positional associations of the pulsars in this catalogwith other astrophysical sources. For the EGRET 3EG, EGR, and GEV and AGILE AGL catalogs,the uncertainties in the localization of the counterparts is worse than for the LAT sources. In thesecases, we consider a source is a possible counterpart to a LAT pulsar when the separation betweenthe two positions is less than the quadratic sum of their 95% confidence error radii.We see that 25 of the 46 pulsars are associated with sources in the 3EG, EGR and GEVcatalogs of EGRET sources, though 19 were seen only as unidentified unpulsed sources. A numberof these unidentified EGRET sources had previously been associated with SNRs, PWNe, or otherobjects (e.g. Walker et al. 2003; De Becker et al. 2005). In all cases, the gamma-ray emission seenwith the LAT is dominated by the pulsed emission. Of the 25 EGRET sources, 14 are gamma-ray-selected pulsars, and 11 are radio-selected, including 2 MSPs. All 6 high-confidence EGRET pulsars(Nolan et al. 1996) are detected, and the 3 marginal EGRET detections are confirmed as pulsars(Ramanamurthy et al. 1996; Kaspi et al. 2000; Kuiper et al. 2000). The 21 sources without 3EG,EGR, or GEV counterparts include 18 previously detected radio pulsars (6 of which are MSPs) and3 gamma-ray selected pulsars. 21 –Not surprisingly, many of the young pulsars have SNR or PWN associations. At least 19 ofthe 46 pulsars are associated with a PWN and/or SNR (Roberts et al. 2005; Green 2009). We donot test here whether the gamma-ray flux from any of these pulsars includes a non-magnetosphericcomponent, as might be indicated by spatially extended emission or a spectrum at pulse minimumnot characteristic of a pulsar. Such studies are underway.At least 12 of the pulsars are associated with TeV sources, 9 of which are also associated withPWNe. Those pulsars with both TeV and PWN associations are typically young, with ages lessthan 20 kyr.
4. Pulsar Flux Sensitivity
In order to interpret the population of gamma-ray pulsars discovered with the LAT, we needto evaluate the sensitivity of our searches for pulsed emission. While the precise sensitivity at anylocation is a function of the local background flux, the pulsar spectrum, and the pulse shape, we canderive an approximate pulsed sensitivity by calculating the unpulsed flux sensitivity for a typicalpulsar spectrum at all locations in the sky and correlating with the observed Z test statistic forthe ensemble of detected pulsars.Figures 7 and 8 show the distributions of the cutoff energy and the photon index, respectively,for all the LAT-detected pulsars. The distribution of photon indices peaks in the range Γ = 1 − E cutoff = 1 − . E cutoff = 2 . (l,b) location in the sky, we computed the DC flux sensitivity at a thresholdlikelihood test statistic TS = 25 integrated above 100 MeV, assuming the typical pulsar spectrumwithin the source PSF and an underlying diffuse gamma-ray flux from the rings Galaxy v0 model(Abdo et al. 2009f). This is an earlier version of the publicly available model gll iem v02 , similarto the model used for the spectral analysis. We note that the likelihood calculation assumes that thesource flux is small compared to the diffuse background flux within the PSF, which is appropriatefor a source just at the detection limit. Finally, we converted this map to pulsed sensitivity by asimple scale factor that accounts for the correspondence between the Z periodicity test confidencelevel and the unpulsed likelihood TS for the detected pulsars.The resulting 5 σ sensitivity map for pulsed emission is shown in Figure 9. Comparing themeasured fluxes with the predicted sensitivities at the pulsar locations (Figure 10), we see that this5 σ limit indeed provides a reasonable lower envelope to the pulsed detections in this catalog. Thusthe effective sensitivity for high latitude (e.g. millisecond) pulsars with known rotation ephemeridesis 1 − × − cm − s − ; at low latitude there is large variation, with typical detection thresholds3 − × higher. We expect the threshold to be somewhat higher for pulsars found in blind periodsearches. Figure 10 suggests that this threshold is 2 − × higher than that for pulsars discovered 22 –in folding searches, with resulting values as high as 2 × − cm − s − on the Galactic plane.The Log N–Log S plot is shown in Figure 11. The approximate N ∝ /S dependence expectedfor a disk population is apparent for the higher flux objects. This shows that while radio-selectedpulsars are detected down to a threshold of 2 × − cm − s − , the faintest gamma-ray-selectedpulsar detected has a flux ∼ × higher at 6 × − cm − s − . It is interesting to note that, asidefrom the lower flux threshold for the former, the radio-selected and gamma-ray-selected histogramsare well matched, suggesting similar underlying populations.
5. Discussion
The striking results of the early
Fermi pulsar discoveries demonstrate the LAT’s excellentpower for pulsed gamma-ray detection. By increasing the gamma-ray pulsar sample size by nearlyan order of magnitude and by firmly establishing the gamma-ray-selected (radio-quiet Geminga-type) and millisecond gamma-ray pulsar populations, we have promoted GeV pulsar astronomy toa major probe of the energetic pulsar population and its magnetospheric physics. Our large pulsarsample allows us both to establish patterns in the pulse emission possibly pointing to a commonorigin of pulsar gamma-rays and to find anomalous systems that may point to exceptional pulsargeometries and/or unusual emission physics. In this Section we discuss some initial conclusionsdrawn from the sample, recognizing that the full exploitation of these new results will flow fromthe detailed population and emission physics studies to follow.
A widely-cited predictor of gamma-ray pulsar detectability is the spin-down flux at Earth ˙
E/d (see e.g. Smith et al. 2008). However, as argued by Arons (2006) (see also Harding & Muslimov2002), it is natural in many models for the gamma-ray emitting gap to maintain a fixed voltagedrop. This implies that L γ is simply proportional to the particle current (Harding 1981), whichgives L γ ∝ ˙ E / , i.e. gamma-ray efficiency increases with decreasing spin-down power down to˙ E ∼ − erg s − where the gap saturates at large efficiency. In Figure 12 we show how ourdetected pulsars rank in ˙ E / /d against the set of searched pulsars. We see that for both MSPsand young pulsars, the detected objects have among the largest values of this metric. The presenceof missing objects among the detected pulsars is interesting, but must be treated with caution,as the detectability metric may be inflated by poor DM distances, or the sensitivity of the pulsesearch might be anomalously low due to high local background or unfavorable pulsar spectrumor pulse profile. Alternatively, some missing objects may be truly gamma-ray faint for the Earthline-of-sight. A more complete study of the implications of the pulsar non-detections and upperlimits is in progress.To study the luminosity evolution in the observed pulsar population, we plot in Figure 6 our 23 –present best estimate of L γ against ˙ E , based on the pulsed flux measured for each pulsar. Twoimportant caveats must be emphasized here. First, the inferred luminosities are quadraticallysensitive to the often large distance uncertainties. Indeed, for many radio selected pulsars (greenpoints) we have only DM-based distance estimates. For many gamma-ray-selected pulsars we haveonly rather tenuous SNR or birth cluster associations with rough distance bounds. Only a handfulof pulsars have secure parallax-based distances. Second, we have assumed here uniform phase-averaged beaming across the sky ( f Ω =1). This is not realized for many emission models, especiallyfor low ˙ E pulsars (Watters et al. 2009).To guide the eye, Figure 6 shows lines for 100% conversion efficiency ( L γ = ˙ E ) and a heuristicconstant voltage line L γ = (10 erg s − ˙ E ) / . In view of the large luminosity uncertainties, wemust conclude that it is not yet possible to test the details of the luminosity evolution. However,some trends are apparent and individual objects highlight possible complicating factors. For thehighest ˙ E pulsars, there does seem to be rough agreement with the ˙ E / trend. However, largevariance between different distance estimates for the Vela-like PSRs J2021+3651 and J1709 − erg s − < ˙ E < . erg s − , the L γ seemsnearly constant, although the lack of precise distance measurements limits our ability to drawconclusions. For example, the very large nominal DM distance of PSR J0248+6021 would require > τ c ∼ × larger than the age of the putative associated SNR γ Cygni. Improveddistance estimates in this range are the key to probing luminosity evolution.From 10 erg s − < ˙ E < erg s − we have several nearby pulsars with reasonably ac-curate parallax distance estimates. However we see a wide range of gamma-ray efficiencies. Thisis the range over which, for both slot gap and outer gap models, the gap is expected to ‘satu-rate’ and use most of the available potential to maintain the pair cascade. In slot gap models(Muslimov & Harding 2003), the break occurs at about 10 erg s − , when the gap is limited byscreening of the accelerating field by pairs. The efficiency below this saturation is predicted tobe ∼ E ∼ erg s − . With the present statistics and uncertainties, it is not possible to discrim-inate between these model predictions except to note that both are consistent with the observedresults. In some models the gap saturation dramatically affects the shape of the beam on the skyand accordingly the flux conversion factor f Ω ; for outer gap models Watters et al. (2009) estimate f Ω ∼ . − .
15 for Geminga (similar values are obtained for J1836+5925), driving down the ratherhigh inferred luminosity of these pulsars by an order of magnitude. In contrast, another pulsarwith an accurate parallax distance, PSR J0659+1414, has an inferred luminosity 30 × lower thanthe ˙ E / prediction. Clearly, some parameter in addition to ˙ E controls the observed L γ . Finally,for < erg s − the sample is dominated by the MSPs. These nearby, low luminosity objects 24 –clearly lie below the ˙ E / trend, and in fact seem more consistent with L γ ∝ ˙ E .Upper limits on radio pulsars with high values of the spin-down flux at Earth or large ˙ E / /d can help constrain viable efficiency models. In practice, the modest present exposure, the largebackground in the Galactic plane and the need to rely on uncertain dispersion-based distanceestimates limit the value of such constraints. Still, a few pulsars are already interesting; for example,using the DM-based distance, the sensitivity in Figure 9, and an assumed f Ω = 1, we find that PSRJ1740+1000 shows less than 1 / E / (constant voltage) line in Figure6. Similarly, PSRs J1357 − | ζ | ∼ ◦ . In outer gap models this makes it highly unlikely that the pulsar will pro-duce strong emission on the Earth line-of-sight. Similarly it has been argued that PSR J0659+1414has a small viewing angle ζ < ◦ (Everett & Weisberg 2001) (but see Weltevrede & Wright (2009)for a discussion of uncertainties). Again, strong emission from above the null charge surface is notexpected for this ζ . One possible interpretation is that we are seeing slot gap or even polar capemission from this pulsar, which is expected at this ζ . The unusual pulse profile and spectrum ofthis pulsar may allow us to test this idea of alternate emission zones.In discussing non-detections, we should also note that the only binary pulsar systems reportedin this paper are the radio-timed MSPs. In particular, our blind searches are not, as yet, sensitiveto pulsars that are undergoing strong acceleration in binary systems. However, we do expect suchobjects to exist. Population syntheses (Pfahl et al. 2002) suggest that several percent of the youngpulsars are born while retained in massive star binary systems. A few such systems are known inthe radio pulsar sample (e.g. the TeV-detected PSR B1259 − ◦
303 (Abdo et al. 2009h) and LS 5039 (Abdo et al. 2009j) may host pulsed GeV signals thathave not yet been found.
With the above caveats about missing binary systems in mind, we can already draw someconclusions about the single gamma-ray pulsar population. For example, there are 17 gamma-rayselected pulsars with a faintest flux of ∼ × − cm − s − ; there are 16 non-millisecond radio-selected pulsars to this flux limit. Of course, some gamma-ray-selected objects can indeed bedetected in the radio (Camilo et al. 2009). Indeed, the detection of PSR J1741 − L . ≈
25 –0 .
03 mJy kpc underlines the fact that the radio emission can be very faint. Deep searches foradditional radio counterparts are underway. However, with deep radio observations of severalobjects, e.g. Geminga, PSR J0007+7303, PSR J1836+5925 (Kassim & Lazio 1999; Halpern et al.2004, 2007), providing no convincing detections, it is clear that some objects are truly radio faint.The substantial number of radio faint objects suggests that gamma-ray emission has an appreciablylarger extent than the radio beams, such as expected in the outer gap (OG) and slot-gap/two polecaustic (SG/TPC) models.Population synthesis studies for normal (non-millisecond) pulsars predicted that LAT woulddetect from 40–80 radio loud pulsars and comparable numbers of radio quiet pulsars in the firstyear (Gonthier et al. 2004; Zhang et al. 2007). The ratio of radio-selected to gamma-ray-selectedgamma-ray pulsars has been noted as a particularly sensitive discriminator of models, since theouter magnetosphere models predict much smaller ratios than polar cap models (Harding et al.2007). Studies of the MSP population (Story et al. 2007) predicted that LAT would detect around12 radio-selected and 33–40 gamma-ray-selected MSPs in the first year, in rough agreement with thenumber of radio-selected MSPs seen to date (searches for gamma-ray selected MSPs have not yetbeen conducted). Thus, in the first six months the numbers of LAT pulsar detections are consistentwith the predicted range, and the large number of gamma-ray selected pulsars discovered so earlyin the mission points towards the outer magnetosphere models.We can in fact use our sample of detected gamma-ray pulsars to estimate the Galactic birthrates.For each object with an available distance estimate, we compute the maximum distance for detec-tion from D max = D est ( F γ /F min ) / , where D est comes from Table 5, the photon flux F fromTable 4 and F min from Figure 9. We limit D max to 15 kpc, and compare V , the volume enclosedwithin the estimated source distance, to V max , that enclosed within the maximum distance, for aGalactic disk with radius 10 kpc and thickness 1 kpc. If we assume a blind search threshold 2 × higher than that for a folding search at a given sky position, the inferred values of h V /V max i are0.49, 0.59 and 0.55 for the radio-selected young pulsars, millisecond pulsars and gamma-ray-selectedpulsars, respectively. These are close to 0.5, the value expected for a population uniformly filling agiven volume (Schmidt 1968); the MSP value is somewhat high as our sample includes three objectsdetected at < σ . The value for the gamma-ray-selected pulsars is also high but is controlled byPSR J2032+4127. If we exclude this object from the sample, we get h V /V max i = 0 . × the ephemeris-folding value.Although we do not attempt a full population synthesis here, if we assume that the pulsarcharacteristic ages are the true ages, our sample can give rough estimates for local volume birthrates:8 × − kpc − yr − (young radio-selected), 4 × − kpc − yr − (young gamma-ray-selected, 2 × threshold) and 2 × − kpc − yr − (MSP). Note that only half of the gamma-ray-selected objectshave distance estimates. If we assume that the set without distance information has comparableluminosity, the gamma-ray-selected birthrate is ∼ × larger. These estimates retain appreciableuncertainty; for example if the effective blind search detection threshold is 3 × that for folding, theinferred gamma-ray-selected birthrate increases by an additional ∼ /
100 yr (radio-selectedyoung pulsars), 1 /
100 yr (gamma-ray-selected pulsars) and 1 / (6 × yr) (radio-selected MSPs).Normally in estimating radio pulsar birthrates one would correct for the radio beaming fraction.However if young gamma-ray-selected pulsars are simply similar objects viewed outside of the radiobeam, this would result in double-counting. In any case one infers a total Galactic birthrate forenergetic pulsars of ∼ /
50 yr, with gamma-ray-selected objects representing half. This represents alarge fraction of the estimated Galactic supernova rate, so clearly more careful population syntheseswill be needed to see if these numbers are compatible.
The pulse shape properties can also help us to probe the geometry and physics of the emissionregion. The great majority of the pulsars show two dominant, relatively sharp peaks, suggestingthat we are seeing caustics from the edge of a hollow cone. When a single peak is seen, it tends tobe broader, suggesting a tangential cut through an emission cone. This picture is realized in theOG and the high altitude portion of the SG models.For the radio-emitting pulsars, we can compare the phase lag between the radio and firstgamma-ray peak δ with the separation of the two gamma-ray peaks ∆. As first pointed out inRomani & Yadigaroglu (1995), these should be correlated in outer magnetosphere models — thisis indeed seen (Figure 4). The distribution can be compared with predictions of the TPC andOG models shown in Watters et al. (2009). The δ − ∆ distribution and in particular the presenceof ∆ ∼ . − . ∼ . − .
5, which favors TPC models. In Figure 13 we show the peak separation as afunction of pulsar spin-down luminosity — the ∆ distribution appears to be bimodal, with no strongdependence on pulsar ˙ E (or age). A full comparison will require detailed population models, whichare being created. It may also be hoped that the precise distribution of measured values can helpprobe details of the emission geometry. In particular, whenever we have external constraints on theviewing angle ζ (typically from X-ray images of the PWN) or magnetic inclination α (occasionallymeasured from radio polarization), then the observed values of δ and ∆ become a powerful probeof the precise location of the emission sheet within the magnetosphere. This can be sensitive to thefield perturbations from magnetospheric currents and hence can probe the global electrodynamicsof the pulsar magnetosphere.If one examines the energy dependence of the light curves of both the radio-selected andgamma-ray-selected pulsars, a decrease in the P1/P2 ratio with increasing energy seems to bea common feature. However, the P1/P2 ratio evolution does not occur for all pulsars, notablyJ0633+0632, J1028 − − − − ∼ B LC are uniformly relatively large( & G). Indeed, the LAT-detected MSPs are those with the highest light cylinder fields withvalues very similar to those of the detected normal pulsars. Comparison of the spectral cut-off E cutoff with surface magnetic field shows no significant correlation. This evidence argues againstclassical low altitude polar cap models supported by γ -B cascades. However, there is a weakcorrelation of E cutoff with B LC , as shown in Figure 7. It is interesting that the values of E cutoff have a range of only about a decade, from 1 to 10 GeV, and that all the different types of pulsarsseem to follow the same correlation. This strongly implies that the gamma-ray emission originatesin similar locations in the magnetosphere relative to the light cylinder. Such a correlation of E cutoff with B LC is actually expected in all outer magnetosphere models where the gamma-ray emissionprimarily comes from curvature radiation of electrons whose acceleration is balanced by radiationlosses. In this case, E cutoff = 0 . λ c (cid:18) E k e (cid:19) / ρ / c (8)in m e c , where λ c is the electron Compton wavelength, E k is the electric field that acceleratesparticles parallel to the magnetic field and ρ c is the magnetic field radius of curvature. In bothSG (Muslimov & Harding 2004) and OG (Zhang et al. 2004; Hirotani 2008) models, E k ∝ B LC w ,where w is the gap width. All these models give values of E cutoff that are roughly consistent withthose measured for the LAT pulsars. Although ρ c ∼ R LC , the gap widths are expected to decreasewith increasing B LC , so that E cutoff is predicted to be only weakly dependent on B LC in most outermagnetosphere models, as observed.The detection of pulsed flux at E max > a few GeV provides additional, physical motivation forhigh altitude emission, since one expects strong (hyper-exponential) attenuation from γ -B −→ e + e − absorption in the high magnetic fields at low altitudes of a near-surface polar gap. For a polarcap model with spin period P and surface field 10 B G, Equation 1 of Baring (2004) gives r & ( E max B / . / P − / cm. While the largest minimum r is derived from the 25 GeVdetection of the Crab by MAGIC (Albert et al. 2008), significant minimum altitudes of 2 to 3 stellarradii are found for many LAT pulsars with large B and high E cutoff , assuming maximum energiesset at around 2 . E cutoff . Such altitudes are inconsistent with the r . − E , although there is a great deal of scatter; a similar trend was noted in Thompson et al. (1999).This may be indicative of higher pair multiplicity, which would steepen the spectrum for the moreenergetic pulsars, either by steepening the spectrum of the curvature radiation-generating primaryelectrons (Romani 1996) or by inclusion of an additional soft spectral component associated withrobust pair formation (Takata & Chang 2007; Harding et al. 2008). In either case, one would expect 28 –steepening from the simple monoenergetic curvature radiation spectrum Γ = 2 / E pulsars. Interestingly, the MSPs do not extend the trend to lower ˙ E . Of course EGRET (andnow the LAT) find strong variations of photon index with phase for the brighter pulsars. A fullunderstanding of photon index trends will doubtless require phase-resolved modeling.
6. Conclusion
The new gamma-ray pulsar populations established by early LAT observations show that weare detecting many nearby young pulsars. In addition we are detecting the millisecond pulsarswith the highest spin-down flux at Earth. Thus we see that the LAT is providing a new, local,but relatively unbiased view of the energetic pulsar population (see Figures 1, 2, and 3). Thesedetections provide a new window into pulsar demographics and physics.We conclude that a large fraction of the local energetic pulsars are GeV emitters. There is alsoa significant correlation with X-ray and TeV bright pulsar wind nebulae. Conversely, we have nowuncovered the pulsar origin of a large fraction of the bright unidentified Galactic EGRET sources,as proposed by several authors (Kaaret & Cottam 1996; Yadigaroglu & Romani 1997). We havealso found plausible pulsar counterparts for several previously detected TeV sources. In this sensethe “mystery” of the unidentified EGRET sources is largely solved. It is possible that the twoLAT-detected massive binaries (LSI +61 ◦ ∼ /
50 yr, a large fraction of the estimated Galacticsupernova rate. Gamma-ray detectable MSPs in the Galactic field are born rarely, ∼ / × yr,but with their long lifetimes are inferred to contribute comparably to the number of (in principle)detectable Galactic gamma-ray pulsars.The data also advance our understanding of emission zone physics. It is now clear that thegamma-ray emission from the brightest pulsars arises largely in the outer magnetosphere. The pho-ton emission also accounts for a large fraction of the spin-down luminosity, increasing as the pulsarsapproach ˙ E ∼ − erg s − . While these wide, bright beams are a boon for population studies, asnoted above, they represent a challenge for theorists trying to understand pulsar magnetospheres.Further LAT pulsar observations and, in particular, the high quality, highly phase-resolved spectranow being obtained for the brightest LAT pulsars will surely sharpen this challenge.The Fermi
LAT Collaboration acknowledges the generous support of a number of agencies 29 –and institutes that have supported the Fermi LAT Collaboration. These include the NationalAeronautics and Space Administration and the Department of Energy in the United States, theCommissariat `a l’Energie Atomique and the Centre National de la Recherche Scientifique / InstitutNational de Physique Nucl´eaire et de Physique des Particules in France, the Agenzia SpazialeItaliana 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 Foundationand the Swedish National Space Board in Sweden.Additional support for science analysis during the operations phase is gratefully acknowledgedfrom the Istituto Nazionale di Astrofisica in Italy and the Centre National d’´Etudes Spatiales inFrance.The Parkes radio telescope is part of the Australia Telescope which is funded by the Com-monwealth Government for operation as a National Facility managed by CSIRO. The Green BankTelescope is operated by the National Radio Astronomy Observatory, a facility of the NationalScience Foundation operated under cooperative agreement by Associated Universities, Inc. TheArecibo Observatory is part of the National Astronomy and Ionosphere Center (NAIC), a nationalresearch center operated by Cornell University under a cooperative agreement with the NationalScience Foundation. The Nan¸cay Radio Observatory is operated by the Paris Observatory, associ-ated with the French Centre National de la Recherche Scientifique (CNRS). The Lovell Telescopeis owned and operated by the University of Manchester as part of the Jodrell Bank Centre forAstrophysics with support from the Science and Technology Facilities Council of the United King-dom. The Westerbork Synthesis Radio Telescope is operated by Netherlands Foundation for RadioAstronomy, ASTRON.
REFERENCES
Abdo, A. A., et al. 2008, Science, 322, 1218, (CTA1)—. 2009a, ApJ, submitted (Geminga)—. 2009b, Science, 325, 848, (Millisecond Pulsars)—. 2009c, Science, 325, 840, (Blind Search Pulsars)—. 2009d, ApJ, 699, L102, (PSR J0205+6449)—. 2009e, ApJ, 695, L72, (PSR J1028 − − ◦ γ -Cygni)—. 2010d, ApJ, 708, 1254, (Crab)Aharonian, F., et al. 2006a, A&A, 456, 245—. 2006b, A&A, 448, L43—. 2006c, A&A, 457, 899—. 2007, A&A, 472, 489—. 2009, A&A, 499, 723Albert, J., et al. 2008, ApJ, 674, 1037Arons, J. 1996, A&AS, 120, C49—. 2006, On the Present and Future of Pulsar Astronomy, 26th meeting of the IAU, Joint Discussion2, 16-17 August, 2006, Prague, Czech Republic, JD02, − − This preprint was prepared with the AAS L A TEX macros v5.2.
38 –Table 1. Measured and intrinsic parameters of LAT-detected pulsars
PSR Type, l b P ˙ P age τ c ˙ E B LC S Ref. ( ◦ ) ( ◦ ) (ms) (10 − ) (kyr) (10 erg s − ) (kG) (mJy)J0007+7303 g a,b < . c,d − × − × e d − × − ×
24 313.1 0.9J0248+6021 r f b − · · · J0437 − d − × − × h − − d − × − × i b − < . h < i − i − d × − × − k − − l − − m − n − − b , < . − i −
60 g b − < . − i − − d × − × · · · J1709 − n − − i − −
31 g b < . − b,t . − d × − × − o − − b − , < . − b < . − b − , < . − o − b < . b,r − < . n b − · · · J2021+3651 r p b,s · · · J2032+4127 g b,t q − − d − × − × m
39 –Table 1—Continued
PSR Type, l b P ˙ P age τ c ˙ E B LC S Ref. ( ◦ ) ( ◦ ) (ms) (10 − ) (kyr) (10 erg s − ) (kG) (mJy)J2238+59 g b · · · Note. — The first two columns are pulsar names and types: r for radio-selected, g for gamma-ray-selected, m for MSPs, and b for binary pulsars. The 3 rd and 4 th columns are Galactic coordinates foreach pulsar. The 5 th and 6 th columns list the period ( P ) and its first derivative ( ˙ P ), corrected for theShklovskii effect (see text). Following are the characteristic age τ c (column 7), the spin-down luminosity˙ E (column 8), and the magnetic field at the light cylinder B LC (column 9) . The last column is theradio flux density at 1400 MHz, or an upper limit when one is available. These values are taken from theATNF database (Manchester et al. 2005) except for the noted entries where: (1) (Halpern et al. 2004); (2)(Camilo et al. 2009); (3) (Roberts et al. 2002); (4) (Halpern et al. 2007); (5) (Ray et al. 1996). Note thatPSR J1509 − −
58 observed by
CGRO .References. — References to
Fermi
LAT publications specific to these pulsars: a (Abdo et al. 2008); b (Abdo et al. 2009c) ; c (Abdo et al. 2009n) ; d (Abdo et al. 2009b) ; e (Abdo et al. 2009d) ; f (Cognard et al. 2010) ; h (Abdo et al. 2010d) ; i (Weltevrede et al. 2010) ; j (Abdo et al. 2009a) ; k (Abdo et al. 2009g) ; l (Abdo et al. 2009e) ; m (Abdo et al. 2009f) ; n (Abdo et al. 2010a) ; o (Camilo et al.2010) ; p (Abdo et al. 2009m) ; q (Noutsos et al. 2010) ; r (Abdo et al. 2010b) ; s (Abdo et al. 2010c) ; t (Camilo et al. 2009).
40 –Table 2. Pulsation detection significances for LAT-detected pulsars
PSR Z value H value maxROI( ◦ ) ObsIDJ0007+7303 2072.1 2371.8 1.0 LJ0030+0451 121.1 362.7 1.0 NJ0205+6449 90.9 206.0 1.0 G, JJ0218+4232 24.7 22.5 1.0 N, WJ0248+6021 57.5 75.1 0.5 NJ0357+32 422.7 450.7 1.0 LJ0437 − − − − − − − − − − −
60 148.2 159.3 1.0 LJ1509 − − − − −
31 141.2 279.6 1.0 LJ1741 − − − − − − − −
41 –Table 2—Continued
PSR Z value H value maxROI( ◦ ) ObsIDNote. — Columns 2 and 3 list the Z (Buccheri et al.1983) and H (de Jager et al. 1989) periodicity test valuesfor E > . σ , with the exceptions of J0218+4232,J0751+1807, J1744 − σ corre-sponds to Z >
36 and
H >
42 ; greater than 7 σ corre-sponds to Z >
61 ; and greater than 10 σ corresponds to Z >
114 (see Section 2.1.3). Column 4 gives the maxi-mum angular radius (maxROI) around the pulsar positionwithin which gamma-ray events were searched for pulsa-tions. The final column indicates the observatories thatprovided ephemerides (see Section 2.1.1 for details): “A”– Arecibo telescope; “G” – Green Bank Telescope; “J” –Lovell telescope at Jodrell Bank; “L” – Large Area Tele-scope; “N” – Nan¸cay Radio Telescope; “P” – Parkes radiotelescope; “W” – Westerbork Synthesis Radio Telescope.
42 –Table 3. Pulse shape parameters of LAT-detected pulsars
PSR Type a Peak Radio lag γ -ray peak separation Off-pulse definitionmultiplicity δ ∆ φ J0007+7303 g 2 ... 0.23 ± ± ± ± ± ± ± ± − ± ± ± − ± ± ± ± ± − ± ± − ± ± − ± ± − ± ± − ± ± − ± ± J1418 − ± − b ± ± −
60 g 2 ... 0.15 ± − b ± ± − ± ± − ± ± − ± −
31 g 2 ... 0.42 ± − ± ± − ± − ± ± − ± − ± − ± − ± ± ± ± ± ± ± ± ± ± ± ± ± ± − ± ±
43 –Table 3—Continued
PSR Type a Peak Radio lag γ -ray peak separation Off-pulse definitionmultiplicity δ ∆ φ J2238+59 g 2 ... 0.50 ± a Types are r=radio-selected, g=gamma-ray-selected, m=millisecond b For some pulse profiles the current dataset does not allow clear discrimination between a single,broad pulse and two unresolved pulses. See the discussion in Weltevrede et al. (2010) regarding PSRsJ1420 − − δ of the first gamma peak from themain radio peak for the radio-detected pulsars (4th column), and the phase difference ∆ between themain gamma-ray peaks (5th column). Column 6 lists the off-pulse phase range used in the spectralanalysis. The boldface entries for PSR J1124-5916 are the corrected values as per an Erratum sentto the ApJ in December 2010. Table 4. Spectral fitting results for LAT-detected pulsars
PSR Type a Photon Flux ( F ) Energy Flux ( G ) Γ E cutoff TS TS cutoff
Luminosity Efficiency b (10 − ph cm − s − ) (10 − erg cm − s − ) (GeV) (10 erg s − ) ( f Ω = 1)J0007+7303 g 30.7 ± ± ± ± ±
38 0.20 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± · · · · · · J0437 − ± ± ± ± ± ± c r 209 ± ± ± ± ±
310 0.001 ± − ± ± ± ± +0 . − . +0 . − . J0631+1036 r 2.8 ± ± ± ± ± ± ± ± · · · · · · J0633+1746 g 305.3 ± ± ± ± +24 − +0 . − . J0659+1414 r 10 ± ± ± ± ± ± − ± ± ± ± +12 − +0 . − . J0751+1807 m 1.35 ± ± ± ± +1 − . +0 . − . J0835 − ± ± ± ± ±
12 0.01 ± − ± ± ± ± ±
73 0.14 ± − ± ± ± ± ±
90 0.08 ± − ± ± ± ± ± ± − ± ± ± ± +34 − +0 . − . J1418 − ± ± ± ± − ± ± ± ± ±
380 0.06 ± −
60 g 17.8 ± ± ± ± · · · · · · J1509 − ± ± ± ± ±
49 0.15 ± − ± ± ± ± ± ± − ± ± ± ± − ± ± ± ± ±
80 0.09 ± −
31 g 25.3 ± ± ± ± · · · · · · J1741 − ± ± ± ± ± ± − ± ± ± ± ± ± − ± ± ± ± − ± ± ± ± ±
140 0.33 ± − ± ± ± ± · · · · · · J1826 − ± ± ± ± · · · · · · Table 4—Continued
PSR Type a Photon Flux ( F ) Energy Flux ( G ) Γ E cutoff TS TS cutoff
Luminosity Efficiency b (10 − ph cm − s − ) (10 − erg cm − s − ) (GeV) (10 erg s − ) ( f Ω = 1)J1833 − ± ± ± ± ±
60 0.01 ± d g 65.6 ± ± ± ± < < ± ± ± ± · · · · · · J1952+3252 r 17.6 ± ± ± ± ±
32 0.02 ± ± ± ± ± · · · · · · J2021+3651 r 67.35 ± ± ± ± +500 − +0 . − . J2021+4026 d g 152.6 ± ± ± ± ±
150 2.2 ± ± ± ± ± ± ± ± ± ± ± − ± ± ± ± +0 . − . +0 . − . J2229+6114 r 32.6 ± ± ± ± ± ± ± ± · · · · · · a Types are r=radio-selected, g=gamma-ray-selected, m=millisecond. b Here, f Ω is assumed to be 1, which can result in an efficiency > c For the Crab the spectral parameters come from Abdo et al. (2010d). d For J1836+5925 and J2021+4026 the spectral parameters come from the phase-averaged analysis (see Section 2.2).Note. — Results of the unbinned maximum likelihood spectral fits for the LAT gamma-ray pulsars (see Section 2.2). Columns 3 and 4 list the on-pulse photonflux F and on-pulse energy flux G respectively. The fits used an exponentially cutoff power-law model (see Eq. 4) with photon index Γ and cutoff energy E cutoff given in columns 5 and 6. The systematic uncertainties on F , G , and Γ due to uncertainties in the Galactic diffuse emission model have been addedin quadrature with the statistical errors. Uncertainties in the instrument response induce additional biases of δF = (+30% , − δG = (+20% , − δ Γ = (+0 . , − . δE cutoff = (+20% , − T S ) for the source significance is provided in column 7. The significance of an exponentialcutoff (as compared to a simple power-law) is indicated by
T S cutoff in column 8, where a value <
10 indicates that the two models are comparable. The totalgamma-ray luminosity L γ and the resulting calculated gamma-ray conversion efficiency η γ ≡ L γ / ˙ E (where f Ω = 1 as described in Section 3.2) are listed incolumns 9 and 10, respectively. The uncertainties in L γ and η include the flux and distance uncertainties. Nevertheless, the strong dependence of these variableson the measured distance (see Table 5) and beaming factor means that they should be considered with care.
46 –Table 5. Pulsar distance estimates
Pulsar Name Distance (kpc) Method a (Ref b )J0007+7303 1.4 ± ± − ± ± − +0 . − . P (18)J0631+1036 0.75–3.62 O (39)J0633+1746 0.250 +0 . − . P (11)J0659+1414 0.288 +0 . − . P (2)J0742 − +1 . − . DM (33)J0751+1807 0.6 +0 . − . P (28)J0835 − +0 . − . P (10)J1028 − ± − ± − ± − +0 . − . O (13)J1418 − − ± − ± − ± − − ± − ± − +0 . − . P (34)J1747 − − ± − ± < ± +2 . − . O (36)J2021+4026 1.5 ± ± − +0 . − . P (18)J2229+6114 0.8–6.5 O (17,20) a K distance evaluation from kinematic model; Pfrom parallax; DM from dispersion measure using theCordes & Lazio (2002) model; O from other measurements.For DM measurements, we assume a minimum distance un-certainty of 30%, as discussed in Section 3.1. b For DM, the reference gives the DM measurement.Note. — The best known distances of 37 pulsars detectedby
Fermi . Nine of the pulsars in the catalog have no dis-
47 – tance estimate and are not included in this table.References. — (1) Bassa et al. (2003); (2) Brisken et al.(2003); (3) Camilo et al. (2009); (4) Camilo et al. (2006);(5) Camilo et al. (2002b); (6) Cognard et al. (2010);(7) Crawford et al. (2006); (8) D’Amico et al. (2001);(9) Deller et al. (2008); (10) Dodson et al. (2003); (11)Faherty et al. (2007); (12) Gaensler et al. (2004); (13)Gonzalez & Safi-Harb (2003); (14) Green & Gull (1982);(15) Greidanus & Strom (1990); (16) Halpern et al. (2007);(17) Halpern et al. (2001a); (18) Hotan et al. (2006);(19) Keith et al. (2008); (20) Kothes et al. (2001); (21)Johnston et al. (1996); (22) Landecker et al. (1980); (23)Lommen et al. (2006); (24) Manchester et al. (2005); (25)Manchester et al. (2001); (26) McGowan et al. (2004); (27)Ng et al. (2005); (28) Nice et al. (2005); (29) Oka et al.(1999); (30) Pineault et al. (1993); (31) Ray et al. (1996);(32) Roberts et al. (1993); (33) Taylor et al. (1993);(34) Toscano et al. (1999); (35) Trimble (1973); (36)Van Etten et al. (2008); (37) Weltevrede & Wright (2009);(38) Yadigaroglu & Romani (1997); (39) Zepka et al.(1996)
48 –Table 6. Positional associations with known GeV and TeV sources for LAT-detected pulsars
PSR Alt. name LAT BSL association a GeV associations b Other associationsJ0007+7303 · · · EGR J0008+7308 PWN G119.5+10.2 GEV J0008+73041AGL J0006+7311J0030+0451 · · · · · ·
J0205+6449 · · · · · · · · ·
SNR/PWN 3C 58 PWN G119.5+10.2 J0218+4232 · · · · · · · · · · · ·
J0248+6021 · · · · · · · · · · · ·
J0357+32 · · · · · · · · ·
J0437 − · · · · · · · · · PWN G253.4 − J0534+2200 Crab 0FGL J0534.6+2201 3EG J0534+2200 SNR/PWN G184.6 − , PSR B0531+21 EGR J0534+2159 HESS J0534+220 GEV J0534+21591AGL J0535+2205J0613 − · · · − · · · · · · J0631+1036 · · · · · · · · ·
J0633+0632 · · · · · ·
EGR J0633+0646GEV J0633+0645J0633+1746 Geminga 0FGL J0634.0+1745 3EG J0633+1751 PWN G195.1+4.3 EGR J0633+1750 MGRO J0632+17 , GEV J0634+17461AGL J0634+1748J0659+1414 PSR B0656+14 · · · · · ·
SNR 203.0+12.0J0742 − − · · · · · · · · · J0751+1807 · · · · · · · · · · · ·
J0835 − − − − , PSR B0833 −
45 EGR J0834 − − GEV J0835 − − − · · · − − · · · GEV J1025 − − −
58 0FGL J1047.6 − − EGR J1048 − − − −
52 0FGL J1058.1 − − · · · EGR J1058 − − − − · · · · · · · · · MSH 11 − , J1418 − · · · − − GEV J1417 − − − − · · · · · · −
49 –Table 6—Continued
PSR Alt. name LAT BSL association a GeV associations b Other associationsEGR J1418 − − GEV J1417 − − − · · · − · · · · · · J1509 − c · · · − − − J1614 − · · · · · · − · · · J1709 − −
44 0FGL J1709.7 − − − , EGR J1710 − − GEV J1709 − − − · · · · · · · · · HESS J1718 − J1732 − · · · − − · · · EGR J1732 − − − · · · − − · · · J1744 − · · · · · · · · · · · · J1747 − · · · · · · − − J1809 − · · · − − − GEV J1809 − − − · · · − − · · · J1826 − · · · − − − GEV J1825 − − − · · · · · · · · · SNR/PWN G21.5 − , HESS J1833 − J1836+5925 · · · · · ·
GEV J1835+59211AGL J1836+5923J1907+06 · · · J1952+3252 d PSR B1951+32 0FGL J1953.2+3249 · · ·
SNR CTB 80 PWN G69.0+2.7 J1958+2846 · · · · · ·
GEV J1957+2859J2021+3651 · · · J2021+4026 · · · γ Cygni J2032+4127 · · · EGR J2033+41171AGL J2032+4102J2043+2740 · · · · · · · · · · · ·
J2124 − · · · − · · · PWN G10.9 − J2229+6114 · · · EGR J2227+6114 MGRO J2228+61 ,
50 –Table 6—Continued
PSR Alt. name LAT BSL association a GeV associations b Other associationsGEV J2227+61011AGL J2231+6109J2238+59 · · · · · · · · · · · · a Source designator from the LAT Bright Source List (Abdo et al. 2009i). b Source designator(s) from the 3rd EGRET (3EG: Hartman et al. 1999), Revised EGRET (EGR:Casandjian & Grenier 2008), High-energy EGRET (GEV Lamb & Macomb 1997) and/or the firstAGILE (1AGL: Pittori et al. 2009) catalogs. c PSR J1509 − −
58 observed by
CGRO (Kuiper et al.1999). d While pulsations from PSR J1952+3252 were detected in EGRET data, it was never catalogedas a point source.Note. — Alternate names for the pulsars in this catalog are given in column 2. Positionalassociations with SNRs, PWNe and selected TeV sources are provided in column 5.References. — 1. Roberts et al. (2005), 2. Green (2009), 3. Abdo et al. (2009k), 4.Aharonian et al. (2006c), 5. Aharonian et al. (2006b), 6. Aharonian et al. (2006a), 7. Hoppe et al.(2009), 8. Aharonian et al. (2007), 9. Djannati-Ata¨ı et al. (2007), 10. Aharonian et al. (2009), 11.Goodman & Sinnis (2009) .
51 –Fig. 1.— Pulsar sky map in Galactic coordinates. Blue squares: gamma-ray-selected pulsars.Red triangles: millisecond gamma-ray pulsars. Green circles: all other radio loud gamma-ray pul-sars. Black dots: Pulsars for which gamma-ray pulsation searches were conducted using rotationalephemerides. Gray dots: Known pulsars which were not searched for pulsations. 52 –
Log [P(s)] −3 −2.5 −2 −1.5 −1 −0.5 0 0.5 1 1.5
Log [ d P / d t] −20−18−16−14−12−10 e r g / s e r g / s e r g / s e r g / s e r g / s e r g / s y y y y y y Fig. 2.— P − ˙ P diagram. Dashed lines: characteristic age τ c . Dot-dashed lines: rotational energyloss rate ˙ E . Blue squares: gamma-ray-selected pulsars. Red triangles: millisecond gamma-raypulsars. Green circles: all other radio loud gamma-ray pulsars. Black dots: Pulsars for whichgamma-ray pulsation searches were conducted using rotational ephemerides. Gray dots: Knownpulsars which were not searched for pulsations. 53 – X [kpc]-14 -12 -10 -8 -6 -4 -2 0 2 4 6 Y [ k p c ] -10-8-6-4-20246810 Fig. 3.— Galactic plane pulsar distribution (polar view). The star represents the Galactic center.The two circles centered at the Earth’s position have radii of 3 kpc and 5 kpc. For pulsars with dif-ferent possible distances, the nearer values from Table 5 are used. Note that the millisecond pulsars(MSPs), while having a significantly lower ˙ E than the other pulsars (see Figure 8), are detectabledue to their close proximity. The one exception (PSR J0218+4232) also exhibits a significantlyhigher ˙ E than the other MSPs. Blue squares: gamma-ray-selected pulsars. Red triangles: millisec-ond gamma-ray pulsars. Green circles: all other radio loud gamma-ray pulsars. Black dots: Pulsarsfor which gamma-ray pulsation searches were conducted using rotational ephemerides. Gray dots:Known pulsars which were not searched for pulsations. 54 –Fig. 4.— Phase difference ∆ between the gamma-ray peaks, versus the phase lag δ between themain radio peak and the nearest gamma-ray peak. Pulsars without a radio detection are plottedwith δ = 0. With present light curves we cannot generally measure ∆ < .
15; objects classified assingle-peaked are plotted with ∆=0. Two such objects, both MSPs, are off the plot at δ > . δ for PSR J1124-5916, as per an Erratum sent to the ApJ (December 2010). 55 – (yr)] c τ Log [2 4 6 8 10 ( G ) ] L C Log [ B -10123456 Crab
PSR J0357+32 PSR J0218+4232
Fig. 5.— Magnetic field strength at the light cylinder B LC versus pulsar characteristic age τ c .Blue squares: gamma-ray-selected pulsars. Red triangles: millisecond gamma-ray pulsars. Greencircles: all other radio loud gamma-ray pulsars. Black dots: Pulsars for which gamma-ray pulsationsearches were conducted using rotational ephemerides. Gray dots: Known pulsars which were notsearched for pulsations. 56 – )] -1 (erg sELog [33 34 35 36 37 38 39 ) ] - ( e r g s ga mm a Log [ L J1836 + * G e m i ng a V e l a C r a b * B - B + B1706-44
J0437 - + + EL= E ∝ L J2021 + * J0659 + + Fig. 6.— Gamma-ray luminosity L γ versus the rotational energy loss rate ˙ E . Dashed line: L γ equal to ˙ E . Dot-dashed line: L γ proportional to the square root of ˙ E . L γ is calculated using abeam correction factor f Ω = 1 for all pulsars and the integral energy flux G from the on-pulsespectral analysis (see Section 2.2), except for PSRs J1836+5925 and J2021+4026 which use the totalbackground-corrected phase-averaged flux, including a relatively bright unpulsed component (seeSection 2.2). For the Crab we also plot the total high energy luminosity, L tot = L X + L γ , indicatedby *. Several notable pulsars have been labeled. Blue squares: gamma-ray-selected pulsars. Redtriangles: millisecond gamma-ray pulsars. Green circles: all other radio loud gamma-ray pulsars.Unfilled markers indicate pulsars for which only a DM-based distance estimate is available (seeTable 5). Pulsars with two distance estimates have two markers connected with dashed error bars. 57 –Fig. 7.— Value of the exponential cutoff E cutoff versus the magnetic field at the light cylinder, B LC .The statistical uncertainties on E cutoff are shown. An additional systematic bias of (+20% , − E cutoff (see text). The histogram of E cutoff values is projected along the right-hand axis.Blue squares: gamma-ray-selected pulsars. Red triangles: millisecond gamma-ray pulsars. Greencircles: all other radio loud gamma-ray pulsars.Fig. 8.— Photon index Γ versus the rotational energy loss rate, ˙ E . For Γ, the statistical uncertain-ties combined with the systematic uncertainties due to the diffuse emission model are shown. Anadditional systematic bias of (+0 . , − .
1) affects Γ (see text). The histogram of the photon indicesis projected along the right-hand axis. Blue squares: gamma-ray-selected pulsars. Red triangles:millisecond gamma-ray pulsars. Green circles: all other radio loud gamma-ray pulsars. 58 –Fig. 9.— Aitoff projection sky map of the 5 σ sensitivity in units of logarithmic photon flux(Log( L γ ) ph cm − s − ) for six months of Fermi
LAT sky-survey data. The sensitivity analysisuses the model of the diffuse gamma-ray background described in the text (Section 4), and pulsarspectra with differential photon indices of Γ = 1 . E cutoff = 2 . F , versus the 5 σ flux sensitivity de-scribed in Figure 9. For F , the statistical uncertainties combined with the systematic uncertain-ties due to the diffuse emission model are shown. An additional systematic bias of (+30% , − F (see text). The effective blind search sensitivity is comparable to the 2 × σ line,although a few pulsars are discovered at lower flux, presumably due to favorable pulse profiles,spectra or local backgrounds. Blue squares: gamma-ray-selected pulsars. Red triangles: millisec-ond gamma-ray pulsars. Green circles: all other radio loud gamma-ray pulsars. 59 – ) -1 s -2 Log Flux (ph cm-8 -7.5 -7 -6.5 -6 -5.5 -5 -4.5 N u m be r w i t h F l u x > F All gamma-ray psrRadio selectedGamma selected
Fig. 11.— Log N–Log S distribution as described in Section 4 for all the detected pulsars (blackdashed line), the radio-selected gamma-ray pulsars including MSPs (grey histogram), and thegamma-ray-selected pulsars (blue hatched histogram). 60 – (cid:0) (cid:1) (cid:2) (cid:3) (cid:4) (cid:5) (cid:6) (cid:7) (cid:8) (cid:9) (cid:10) L og [˙ E / D ] ( s c a l e d t o V e l a ) Fig. 12.— Pulsar “detectability” metric ( ˙ E / /d , normalized to Vela) vs. spin period. DetectedMSPs (red triangles) and young pulsars (radio-selected, green circles; gamma-ray-selected, bluesquares) all have high values of this metric. For objects with a distance range in Table 5, we usehere the geometric mean of the maximum and minimum values. Searched, but presently undetectedobjects (gray dots) are plotted using DM-derived distances. For the possible causes of non-detectionsee Section 5.1. 61 –Fig. 13.— Separations ∆ between the gamma-ray peaks, for those pulsars with two identifiedpeaks, versus the spin-down power ˙ E . The histogram of peak separations is projected along theright-hand axis. Blue squares: gamma-ray-selected pulsars. Red triangles: millisecond gamma-raypulsars. Green circles: all other radio loud gamma-ray pulsars. 62 – Appendix: Gamma-ray Pulsar Light Curves C oun t s >0.1 GeV J0007+7303 C oun t s >1.0 GeV C oun t s Pulsar Phase C oun t s Fig. A-1.— Light curves for PSR J0007+7303 ( P =316 ms). C oun t s >0.1 GeV J0030+0451 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) −0.200.20.40.60.81 Nancay 1.4 GHz
Fig. A-2.— Light curves for PSR J0030+0451 ( P =4 .
87 ms). C oun t s >0.1 GeV J0205+6449 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) −0.200.20.40.60.81 GBT 2 GHz
Fig. A-3.— Light curves for PSR J0205+6449 ( P =65 . C oun t s >0.1 GeV J0218+4232 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) −0.200.20.40.60.81 Nancay 1.4 GHz
Fig. A-4.— Light curves for PSR J0218+4232 ( P =2 .
32 ms). C oun t s >0.1 GeV J0248+6021 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) Nancay 1.4 GHz
Fig. A-5.— Light curves for PSR J0248+6021 ( P =217 ms). C oun t s >0.1 GeV J0357+32 C oun t s >1.0 GeV C oun t s Pulsar Phase C oun t s Fig. A-6.— Light curves for PSR J0357+32 ( P = 444 ms). C oun t s >0.1 GeV J0437−4715 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) Parkes 3 GHz
Fig. A-7.— Light curves for PSR J0437 − P =5 .
76 ms). C oun t s >0.1 GeV J0534+2200 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) Nancay 1.4 GHz
Fig. A-8.— Light curves for PSR J0534+2200 ( P =33 . C oun t s >0.1 GeV J0613−0200 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) −0.200.20.40.60.81 Nancay 1.4 GHz
Fig. A-9.— Light curves for PSR J0613 − P =3 .
06 ms). C oun t s >0.1 GeV J0631+1036 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) Nancay 1.4 GHz
Fig. A-10.— Light curves for PSR J0631+1036 ( P =288 ms). C oun t s >0.1 GeV J0633+0632 C oun t s >1.0 GeV C oun t s Pulsar Phase C oun t s Fig. A-11.— Light curves for PSR J0633+0632 ( P =297 ms). C oun t s >0.1 GeV J0633+1746 C oun t s >1.0 GeV C oun t s Pulsar Phase C oun t s Fig. A-12.— Light curves for PSR J0633+1746 ( P =237 ms, Geminga pulsar). C oun t s >0.1 GeV J0659+1414 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) Nancay 1.4 GHz
Fig. A-13.— Light curves for PSR J0659+1414 ( P =385 ms). C oun t s >0.1 GeV J0742−2822 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) Nancay 1.4 GHz
Fig. A-14.— Light curves for PSR J0742 − P =167 ms). C oun t s >0.1 GeV J0751+1807 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) Nancay 1.4 GHz
Fig. A-15.— Light curves for PSR J0751+1807 ( P =3 .
48 ms). C oun t s >0.1 GeV J0835−4510 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) Parkes 1.4 GHz
Fig. A-16.— Light curves for PSR J0835 − P =89 . C oun t s >0.1 GeV J1028−5820 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) Parkes 1.4 GHz
Fig. A-17.— Light curves for PSR J1028 − P =91 . C oun t s >0.1 GeV J1048−5832 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) Parkes 1.4 GHz
Fig. A-18.— Light curves for PSR J1048 − P =124 ms). C oun t s >0.1 GeV J1057−5226 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) Parkes 1.4 GHz
Fig. A-19.— Light curves for PSR J1057 − P =197 ms, PSR B1055 − C oun t s > 0.1 GeV J1124-5916 C oun t s > 1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) -0.200.20.40.60.81 Parkes 1.4 GHz
Fig. A-20.— Light curves for PSR J1124 − P =135 ms). The radio phase alignment has been corrected,as per an Erratum sent to the ApJ (December 2010). C oun t s >0.1 GeV J1418−6058 C oun t s >1.0 GeV C oun t s Pulsar Phase C oun t s Fig. A-21.— Light curves for PSR J1418 − P =111 ms). C oun t s >0.1 GeV J1420−6048 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) −0.200.20.40.60.81 Parkes 1.4 GHz
Fig. A-22.— Light curves for PSR J1420 − P =68 . C oun t s >0.1 GeV J1459−60 C oun t s >1.0 GeV C oun t s Pulsar Phase C oun t s Fig. A-23.— Light curves for PSR J1459 −
60 ( P =103 ms). C oun t s >0.1 GeV J1509−5850 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) −0.4−0.200.20.40.60.81 Parkes 1.4 GHz
Fig. A-24.— Light curves for PSR J1509 − P =88 . C oun t s >0.1 GeV J1614−2230 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) GBT 2 GHz
Fig. A-25.— Light curves for PSR J1614 − P =3 .
15 ms). C oun t s >0.1 GeV J1709−4429 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) Parkes 1.4 GHz
Fig. A-26.— Light curves for PSR J1709 − P =102 ms, PSR B1706 − C oun t s >0.1 GeV J1718−3825 C oun t s >1.0 GeV C oun t s C oun t s Rotational Phase R a d i o F l u x ( a u ) −0.200.20.40.60.81 Parkes 1.4 GHz
Fig. A-27.— Light curves for PSR J1718 − P =74 . C oun t s >0.1 GeV J1732−31 C oun t s >1.0 GeV C oun t s Pulsar Phase C oun t s Fig. A-28.— Light curves for PSR J1732 −
31 ( P =197 ms). C oun t s >0.1 GeV J1741−2054 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) −0.200.20.40.60.81 GBT 2 GHz
Fig. A-29.— Light curves for PSR J1741 − P =414 ms). While this pulsar is detected in the radio(Camilo et al. 2009), it was discovered by the LAT andis considered a gamma-ray-selected pulsar. C oun t s >0.1 GeV J1744−1134 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) Nancay 1.4 GHz
Fig. A-30.— Light curves for PSR J1744 − P =4 .
08 ms). C oun t s >0.1 GeV J1747−2958 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) −0.200.20.40.60.81 GBT 2 GHz
Fig. A-31.— Light curves for PSR J1747 − P =98 . C oun t s >0.1 GeV J1809−2332 C oun t s >1.0 GeV C oun t s Pulsar Phase C oun t s Fig. A-32.— Light curves for PSR J1809 − P =147 ms). C oun t s >0.1 GeV J1813−1246 C oun t s >1.0 GeV C oun t s Pulsar Phase C oun t s Fig. A-33.— Light curves for PSR J1813 − P =48 . C oun t s >0.1 GeV J1826−1256 C oun t s >1.0 GeV C oun t s Pulsar Phase C oun t s Fig. A-34.— Light curves for PSR J1826 − P =110 ms). C oun t s >0.1 GeV J1833−1034 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) GBT 0.8 GHz
Fig. A-35.— Light curves for PSR J1833 − P =61 . C oun t s >0.1 GeV J1836+5925 C oun t s >1.0 GeV C oun t s Pulsar Phase C oun t s Fig. A-36.— Light curves for PSR J1836+5925 ( P =173 ms). C oun t s >0.1 GeV J1907+06 C oun t s >1.0 GeV C oun t s Pulsar Phase C oun t s Fig. A-37.— Light curves for PSR J1907+06 ( P =107 ms). C oun t s >0.1 GeV J1952+3252 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) −0.200.20.40.60.81 Nancay 1.4 GHz
Fig. A-38.— Light curves for PSR J1952+3252 ( P =39 . C oun t s >0.1 GeV J1958+2846 C oun t s >1.0 GeV C oun t s Pulsar Phase C oun t s Fig. A-39.— Light curves for PSR J1958+2846 ( P =290 ms). C oun t s >0.1 GeV J2021+3651 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) GBT 2 GHz
Fig. A-40.— Light curves for PSR J2021+3651 ( P =104 ms). C oun t s >0.1 GeV J2021+4026 C oun t s >1.0 GeV C oun t s Pulsar Phase C oun t s Fig. A-41.— Light curves for PSR J2021+4026 ( P =265 ms). C oun t s >0.1 GeV J2032+4127 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) GBT 2 GHz
Fig. A-42.— Light curves for PSR J2032+4127 ( P =143 ms). While this pulsar is detected in the radio(Camilo et al. 2009), it was discovered by the LAT andis considered a gamma-ray-selected pulsar. C oun t s >0.1 GeV J2043+2740 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) Nancay 1.4 GHz
Fig. A-43.— Light curves for PSR J2043+2740 ( P =96 . C oun t s >0.1 GeV J2124−3358 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) −0.200.20.40.60.81 Nancay 1.4 GHz
Fig. A-44.— Light curves for PSR J2124 − P =4 .
93 ms). C oun t s >0.1 GeV J2229+6114 C oun t s >1.0 GeV C oun t s C oun t s Pulsar Phase R a d i o F l u x ( a u ) −0.200.20.40.60.81 GBT 2 GHz