On 2017 August 17 the gravitationalwave event GW170817 was observed by the Advanced LIGO and Virgo detectors and the gammaray burst GRB GRB 170817A was observed independently by the Fermi Gammaray Burst Monitor and the AntiCoincidence Shield for the Spectrometer for the International GammaRay Astrophysics Laboratory. The probability of the nearsimultaneous temporal and spatial observation of GRB 170817A and GW170817 occurring by chance is 5.0 10SUP8SUP. We therefore confirm binary neutron star mergers as a progenitor of short GRBs. The association of GW170817 and GRB 170817A provides new insight into fundamental physics and the origin of short GRBs. We use the observed time delay of 1.74 0.05 s between GRB 170817A and GW170817 to I constrain the difference between the speed of gravity and the speed of light to be between 3 10SUP15SUP and 7 10SUP16SUP times the speed of light II place new bounds on the violation of Lorentz invariance III present a new test of the equivalence principle by constraining the Shapiro delay between gravitational and electromagnetic radiation. We also use the time delay to constrain the size and bulk Lorentz factor of the region emitting the gammarays. GRB 170817A is the closest short GRB with a known distance but is between 2 and 6 orders of magnitude less energetic than other bursts with measured redshift. A new generation of gammaray detectors and subthreshold searches in existing detectors will be essential to detect similar short bursts at greater distances. Finally we predict a joint detection rate for the Fermi Gammaray Burst Monitor and the Advanced LIGO and Virgo detectors of 0.11.4 per year during the 20182019 observing run and 0.31.7 per year at design sensitivity.
['binaries: close', 'gamma-ray burst: general', 'gravitational waves', 'Astrophysics - High Energy Astrophysical Phenomena', 'General Relativity and Quantum Cosmology']
Gravitational Waves and GammaRays from a Binary Neutron Star Merger GW170817 and GRB 170817A
{'No Header': '<!-- image -->', 'Gravitational Waves and Gamma-Rays from a Binary Neutron Star Merger: GW170817 and GRB 170817A': 'LIGO Scienti /uniFB01 c Collaboration and Virgo Collaboration, Fermi Gamma-ray Burst Monitor, and INTEGRAL ( See the end matter for the full list of authors. )  \nReceived 2017 October 6; revised 2017 October 9; accepted 2017 October 9; published 2017 October 16', 'Abstract': "On 2017 August 17, the gravitational-wave event GW170817 was observed by the Advanced LIGO and Virgo detectors, and the gamma-ray burst ( GRB ) GRB /uniF0A0 170817A was observed independently by the Fermi Gamma-ray Burst Monitor, and the Anti-Coincidence Shield for the Spectrometer for the International Gamma-Ray Astrophysics Laboratory . The probability of the near-simultaneous temporal and spatial observation of GRB /uniF0A0 170817A and GW170817 occurring by chance is 5.0 10 8 ' -. We therefore con /uniFB01 rm binary neutron star mergers as a progenitor of short GRBs. The association of GW170817 and GRB /uniF0A0 170817A provides new insight into fundamental physics and the origin of short GRBs. We use the observed time delay of 1.74 0.05 s +  ( ) between GRB /uniF0A0 170817A and GW170817 to: ( i ) constrain the difference between the speed of gravity and the speed of light to be between 3 10 15 - ' -and 7 10 16 +' -times the speed of light, ( ii ) place new bounds on the violation of Lorentz invariance, ( iii ) present a new test of the equivalence principle by constraining the Shapiro delay between gravitational and electromagnetic radiation. We also use the time delay to constrain the size and bulk Lorentz factor of the region emitting the gamma-rays. GRB /uniF0A0 170817A is the closest short GRB with a known distance, but is between 2 and 6 orders of magnitude less energetic than other bursts with measured redshift. A new generation of gamma-ray detectors, and subthreshold searches in existing detectors, will be essential to detect similar short bursts at greater distances. Finally, we predict a joint detection rate for the Fermi Gamma-ray Burst Monitor and the Advanced LIGO and Virgo detectors of 0.1 -1.4 per year during the 2018 -2019 observing run and 0.3 -1.7 per year at design sensitivity.  \nKey words: binaries: close -gamma-ray burst: general -gravitational waves", '1. Introduction and Background': 'GW170817 and GRB /uniF0A0 170817A mark the discovery of a binary neutron star ( BNS ) merger detected both as a gravitational wave ( GW; LIGO Scienti /uniFB01 c Collaboration & Virgo Collaboration 2017a ) and a short-duration gamma-ray burst ( SGRB; Goldstein et al. 2017; Savchenko et al. 2017b ) . Detecting GW radiation from the coalescence of BNS and neutron star ( NS ) -black hole ( BH ) binary systems has been a major goal ( Abbott et al. 2017a ) of the LIGO ( Aasi et al. 2015 ) and Virgo ( Acernese et al. 2015 ) experiments. This was at least partly motivated by their promise of being the most likely sources of simultaneously detectable GW and electromagnetic ( EM ) radiation from the same source. This is important as joint detections enable a wealth of science unavailable from either messenger alone /uniF0A0 ( Abbott et al. 2017f ) . BNS mergers are predicted to yield signatures across the EM spectrum /uniF0A0 ( Metzger & Berger 2012; Piran et al. 2013 ) , including SGRBs ( Blinnikov et al. 1984; Paczynski 1986;Eichler et al. 1989; Paczynski 1991; Narayan et al. 1992 ) , which produce prompt emission in gamma-rays and longer-lived afterglows.  \nA major astrophysical implication of a joint detection of an SGRB and of GWs from a BNS merger is the con /uniFB01 rmation that these binaries are indeed the progenitors of at least some SGRBs. GRBs are classi /uniFB01 ed as short or long depending on the duration of their prompt gamma-ray emission. This cut is based on spectral differences in gamma-rays and the bimodality of the observed  \n<!-- image -->  \nOriginal content from this work may be used under the terms of the Creative Commons Attribution 3.0 licence. Any further distribution of this work must maintain attribution to the author ( s ) and the title of the work, journal citation and DOI.  \ndistribution of these durations ( Dezalay et al. 1992; Kouveliotou et al. 1993 ) . This empirical division was accompanied by hypotheses that the two classes have different progenitors. Long GRBs have been /uniFB01 rmly connected to the collapse of massive stars through the detection of associated Type Ibc core-collapse supernovae ( see Galama et al. 1998,aswellasHjorth&Bloom 2012 and references therein ) . Prior to the results reported here, support for the connection between SGRBs and mergers of BNSs ( or NS -BH binaries ) came only from indirect observational evidence /uniF0A0 ( Nakar 2007;Bergeretal.2013;Tanviretal.2013; Berger 2014 ) , population synthesis studies ( Bloom et al. 1999; Fryer et al. 1999; Belczynski et al. 2006 ) , and numerical simulations ( e.g., Aloy et al. 2005; Rezzolla et al. 2011;Kiuchi et al. 2015; Baiotti & Rezzolla 2017; Kawamura et al. 2016;Ruiz et al. 2016 ) . The unambiguous joint detection of GW and EM radiation from the same event con /uniFB01 rmsthatBNSmergersare progenitors of ( at least some ) SGRBs.  \nIn Section 2 we describe the independent observations of GW170817 by the LIGO -Virgo and of GRB /uniF0A0 170817A by the Fermi Gamma-ray Burst Monitor ( GBM ) and by the SPectrometer on board INTEGRAL Anti-Coincidence Shield ( SPIACS ) . In Section 3 we establish the /uniFB01 rm association between GW170817 and GRB /uniF0A0 170817A. In Section 4 we explore the constraints on fundamental physics that can be obtained from the time separation between the GW and EM signals. In Section 5 we explore the implications of the joint detection of GW170817 and GRB /uniF0A0 170817A on the SGRB engine and the NS equation of state ( EOS ) . In Section 6 we explore the implications of the comparative dimness of GRB /uniF0A0 170817A relative to the known SGRB population and revise the  \n<!-- image -->  \nFigure 1. Final localizations. The 90% contour for the /uniFB01 nal sky-localization map from LIGO -Virgo is shown in green ( LIGO Scienti /uniFB01 c Collaboration & Virgo Collaboration 2017a, 2017b, 2017c ) . The 90% GBM targeted search localization is overlaid in purple ( Goldstein et al. 2017 ) . The 90% annulus determined with Fermi and INTEGRAL timing information is shaded in gray ( Svinkin et al. 2017 ) . The zoomed inset also shows the position of the optical transient marked as a yellow star ( Abbott et al. 2017f; Coulter et al. 2017a, 2017b ) . The axes are R.A. and decl. in the Equatorial coordinate system.  \n<!-- image -->  \nexpectation rates for joint BNS -SGRB detections in the light of this discovery.', '2. Observational Results': 'The observations of GW170817 and of GRB /uniF0A0 170817A are described in detail in Abbott et al. ( 2017e ) , Goldstein et al. ( 2017 ) , and Savchenko et al. ( 2017b ) . Here we summarize the observations relevant to the results presented in this Letter and report the results of two fully coherent searches for GWs from the sky location of GRB /uniF0A0 170817A. For convenience, all measurements of time have been converted to their geocentric equivalent.', '2.1. LIGO -Virgo Observation of GW170817': "GW170817 is a GW signal from the inspiral of two low-mass compact objects and is the /uniFB01 rst GW observation consistent with a BNS coalescence ( Abbott et al. 2017e, 2017f ) . GW170817 was /uniFB01 rst observed by a low-latency search /uniF0A0 ( Cannon et al. 2012; Messick et al. 2017 ) on 2017 August 17 at 12:41:04 UTC as a single-detector trigger in the LIGO-Hanford detector /uniF0A0 ( Abbott et al. 2017e; LIGO Scienti /uniFB01 c Collaboration & Virgo Collaboration 2017a ) . The temporal proximity of GRB /uniF0A0 170817A was immediately identi /uniFB01 ed by automatic comparison of the Fermi -GBM Gamma-ray Coordinates Network notice to the GW trigger /uniF0A0 ( Urban 2016 ) . Rapid of /uniFB02 ine re-analysis /uniF0A0 ( Usman et al. 2016;Nitzetal.2017b ) of data from the LIGO / Virgo network con /uniFB01 rmed the presence of a signi /uniFB01 cant coincident signal in the LIGO GW detectors with a combined signal-to-noise ratio ( S / N ) of 32.4. The combination of observations from the LIGO and Virgo detectors allowed a precise sky position localization to an area of 28deg 2 at 90% probability shown in green in Figure 1 ( Abbott et al. 2017e; LIGO Scienti /uniFB01 c Collaboration & Virgo Collaboration 2017b ) . A time-frequency representation of the LIGO data containing GW170817 is shown in the bottom panel of Figure 2. The GPS time of the merger of GW170817 is T 1187008882.430 0 GW 0.002 0.002 = -+ s /uniF0A0 ( Abbott et al. 2017e ) .Atthe observed signal strength, the false alarm rate of the all-sky search  \nfor compact-object mergers is less than 1 in 80,000 years ( Abbott et al. 2017e ) . The of /uniFB02 ine searches target binaries with ( detector frame ) total mass 2 -M 500  . Signals are required to be coincident in time and mass in the LIGO detectors, but Virgo data are not used in the signi /uniFB01 cance estimates of the all-sky of /uniFB02 ine search /uniF0A0 ( Abbott et al. 2017e ) .  \nWe present the results of two of /uniFB02 ine targeted searches that coherently combine the data from the LIGO and Virgo detectors and restrict the signal offset time and sky-location using information from the EM observation of GRB /uniF0A0 170817A. The onset of gamma-ray emission from a BNS merger progenitor is predicted to be within a few seconds after the merger, given that the central engine is expected to form within a few seconds and that the jet propagation delays are at most of the order of the SGRB duration ( see, e.g., Finn et al. 1999; Abadie et al. 2012 and references therein ) . The gravitational and EM waves are expected to travel at the same speed.  \nThe /uniFB01 rst targeted search ( Harry & Fairhurst 2011; Williamson et al. 2014; Abbott et al. 2017b;Nitzetal.2017a ) assumes that the source is a BNS or NS -BH binary merger and is located at the sky-position observed for the optical counterpart to GW170817 and GRB /uniF0A0 170817A ( Coulter et al. 2017a, 2017b; Abbott et al. 2017f ) and that there is a 1, 5 - + [ ] s time delay in the arrival of gamma-rays ( determined by the GBM trigger time ) compared to the binary merger time /uniF0A0 ( Abbott et al. 2017b ) .Atthe detection statistic value assigned to GW170817, this search has a p -value of 9.4 10 4.2 6 s < ' > -( ) ,withthissigni /uniFB01 cance estimate limited by computational resources used to estimate the noise background. The second coherent search does not assume any particular GW morphology or GRB model ( Sutton et al. 2010; Was et al. 2012; Abbott et al. 2017b ) and uses the GBM localization of GRB /uniF0A0 170817A to constrain the sky location of the source. This search allows for a 60, 600 -+ [ ] s coincidence between the gamma-rays and the GWs in order to include potentially larger delays in collapsar models of long GRBs. At the detection-statistic value observed for GW170817, this search has a p -value of 1.3 10 4.2 5 s ' -( ) .  \nFigure 2. Joint, multi-messenger detection of GW170817 and GRB /uniF0A0 170817A. Top: the summed GBM lightcurve for sodium iodide ( NaI ) detectors 1, 2, and 5 for GRB /uniF0A0 170817A between 10 and 50 keV, matching the 100 ms time bins of the SPI-ACS data. The background estimate from Goldstein et al. ( 2016 ) is overlaid in red. Second: the same as the top panel but in the 50 -300 keV energy range. Third: the SPI-ACS lightcurve with the energy range starting approximately at 100 keV and with a high energy limit of least 80 MeV. Bottom: the time-frequency map of GW170817 was obtained by coherently combining LIGO-Hanford and LIGOLivingston data. All times here are referenced to the GW170817 trigger time T 0 GW .  \n<!-- image -->  \nThe 90% credible intervals /uniF0A0 ( Veitch et al. 2015;Abbottetal. 2017e ) for the component masses ( in the m m 1 2  convention ) are m M 1.36, 2.26 1 Î  ( ) and m M 0.86, 1.36 2 Î  ( ) , with total mass M 2.82 0.09 0.47 -+  , when considering dimensionless spins with  \nmagnitudes up to 0.89 ( high-spin prior, hereafter ) . When the dimensionless spin prior is restricted to 0.05  ( low-spin prior, hereafter ) , the measured component masses are m 1.36, 1 Î ( M 1.60  ) and m M 1.17, 1.36 2 Î  ( ) , and the total mass is  \nM 2.74 0.01 0.04 -+  . This result is consistent with the masses of all known BNS systems ( Ozel & Freire 2016;Taurisetal.2017 ) . From the GW signal, the best measured combination of the masses is the chirp mass mm m m 1 2 3 5 1 2 1 5  = + ( ) ( ) ,which in the detector frame is found to be M 1.1977 0.0003 0.0008 -+  .  \nThe detection of GW170817 triggered a campaign of EM follow-up observations which led to the identi /uniFB01 cation of NGC 4993 as the host galaxy of GW170817 / GRB /uniF0A0 170817A ( Coulter et al. 2017a, 2017b; Abbott et al. 2017f ) . We evaluate the distance to the host galaxy from the ratio of the Hubble /uniFB02 ow velocity of the host 3017 166 km s 1  -( Abbott et al. 2017g ) and two current measurements of the Hubble constant ( Ade et al. 2016; Riess et al. 2016 ) . These two distance measures are within a combined range of 42.9 3.2  ( ) Mpc, which is consistent with the distance of 40 14 8 -+ Mpc determined with GW data alone and makes GW170817 the closest GW event ever observed ( Abbott et al. 2016a, 2016b, 2017c, 2017d, 2017e ) .  \nThe GW data constrain the inclination angle JN q between the total angular momentum of the system and the line of sight to be anti-aligned, with cos 0.54 JN  q -( Abbott et al. 2017e ) .Asthe binary system component masses are comparable, the NS spins have little impact on the total angular momentum which is aligned with the orbital angular momentum within a few degrees. For discussions in this Letter we will assume that the orbital and total angular momenta are aligned. The SGRB jet is expected to be perpendicular to the accretion disk of the central engine if poweredbyneutrinoannihilation or aligned with the magnetic pole of the rotating central object /uniF0A0 ( Shibata et al. 2006 ) ,hencewe assume the SGRB jet is aligned with the system rotation axis. This yields a jet viewing angle min , 180 56 JN JN  z q q =  - ( ) .As the distance measurement is correlated with JN q , the known distance to NGC 4993 further constrains the viewing angle to  z 36  or  z 28  depending on the assumed value of the Hubble constant ( Abbott et al. 2017g ) , with smaller values of the Hubble constant giving smaller misalignment angles.", '2.2. Fermi-GBM Observation of GRB /uniF0A0 170817A': "GRB /uniF0A0 170817A was autonomously detected in-orbit by the GBM /uniFB02 ight software in a process known as ' triggering. ' Goldstein et al. ( 2017 ) showed the signal exceeds 5 σ in three ( of twelve ) GBM NaI detectors. The GBM detection showed two apparently distinct components. The triggering pulse, that lasts about half a second and falls within the usual observer distributions for GBM SGRBs, is shorter and spectrally harder than the subsequent softer, weaker emission that lasts a few seconds ( Goldstein et al. 2017 ) . Summed GBM lightcurves from the relevant detectors in two energy ranges, selected to show the two distinct components, are shown in the top two panels in Figure 2. The GBM time-tagged event data is binned to match the SPI-ACS temporal resolution ( 100 ms ) and phase ( matching bin edges ) to allow for an easier comparison between the gamma-ray instruments.  \nGoldstein et al. ( 2017 ) quantify the likelihood of GRB /uniF0A0 170817A being an SGRB based only on gamma-ray data. This is done by comparing the measured gamma-ray properties of GRB /uniF0A0 170817A to the known distributions of short and long GRBs. Both the duration distribution alone and the duration and spectral hardness distributions together show that GRB /uniF0A0 170817A is three times more likely to be an SGRB than a long GRB. These analyses are performed in a standard manner, resulting in a longer duration measure than apparent from the  \nhard spike alone because the softer, weaker tail contributes to the calculated duration.  \nThe /uniFB01 nal GBM localization of GRB /uniF0A0 170817A ( including systematic error ) calculated by the GBM targeted search pipeline is shown in Figure 1. This pipeline performs a coherent search over all GBM detectors ( NaI and BGO ) and was originally developed to /uniFB01 nd gamma-ray signals below the onboard triggering threshold around GW triggers ( Blackburn et al. 2015; Connaughton et al. 2016; Goldstein et al. 2016 ) . The 50% and 90% credible regions cover ∼ 350 deg 2 and ∼ 1100 deg 2 , respectively.  \nFitting the main pulse in the GBM data with a parameterized function commonly used for GRB pulses indicates a gammaray emission onset of 0.310 /uniF0A0 ± /uniF0A0 0.048 s before T 0 GBM , where T 0 GBM is de /uniFB01 ned as the time of the GBM trigger ( Goldstein et al. 2017 ) . Based on the position of the optical transient, the signal arrives at Fermi 3.176 ms before it arrives at geocenter. With this correction we /uniFB01 nd that the start of the gamma-ray emission relative to the T 0 GW is 1.74 0.05 s +  ( ) . In this Letter all derived gamma-ray results use 68% con /uniFB01 dence levels.  \nThe spectral analysis using the standard GBM catalog criteria uses data from the 256 ms time interval between T 0.192 s 0 GBM -and T 0.064 s 0 GBM + .A /uniFB01 ttothe ' Comptonized ' function, a power law with a high-energy exponential cutoff ( see Goldstein et al. 2017 for a detailed explanation of this function ) , is preferred over both a simple power-law /uniFB01 t or models with more parameters. The /uniFB01 t produces values of E peak /uniF0A0 = /uniF0A0 ( 215 /uniF0A0 ± /uniF0A0 54 ) keV, and a poorly constrained power-law index 0.14 0.59 a =  . The average /uniFB02 ux for this interval in the 10 -1000 keV range is 5.5 1.2 10 7  ' -( ) erg s -1 cm -2 with a corresponding /uniFB02 uence of 1.4 0.3 10 7  ' -( ) erg cm -2 . The shorter peak interval selection from T 0.128 s 0 GBM -to T 0.064 s 0 GBM -/uniFB01 t prefers the Comptonized function, yielding consistent parameters E 229 78 peak =  ( ) keV, 0.85 1.38 a =  , and peak energy /uniFB02 ux in the 10 -1000 keV of 7.3 2.5 10 7  ' -( ) erg s -1 cm -2 . These standard /uniFB01 tsareusedtocompareGRB /uniF0A0 170817A to the rest of the SGRBs detected by GBM and to place GRB /uniF0A0 170817A in context with the population of SGRBs with known redshift.  \nMore detailed analysis included spectral /uniFB01 ts to the two apparently distinct components. The main emission episode, represented by the peak in Figure 2, appears as a typical SGRB best /uniFB01 t by a power law with an exponential cutoff with spectral index 0.62 0.40 a =- and E 185 62 peak =  ( ) keV over a time interval T 0.320 s 0 GBM -to T 0.256 s 0 GBM + . The time-averaged /uniFB02 ux is 3.1 0.7 10 7  ' -( ) erg s -1 cm -2 . The tail emission that appears spectrally soft is best /uniFB01 tby a blackbody ( BB ) spectrum, with temperature of k T 10.3 1.5 B =  ( ) keV and a time-averaged /uniFB02 ux of 0.53 0.10 10 7  ' -( ) erg s -1 cm -2 , with selected source interval T 0.832 s 0 GBM + to T 1.984 s 0 GBM + . However, this emission is too weak and near the lower energy detection bound of GBM to completely rule out a non-thermal spectrum.  \nThe temporal analysis yielded a T 90,de /uniFB01 ned as the time interval over which 90% of the burst /uniFB02 uence between 50 -300 keV is accumulated, of 2.0 0.5  ( ) s starting at T 0.192 s 0 GBM -. The duration extends beyond the main emission pulse due to the soft component. This analysis reports a 64 ms peak photon /uniFB02 ux of 3.7 0.9  ( ) photons s -1 cm -2 and occurs from T 0.0 s 0 GBM + to T 0.064 s 0 GBM + . The  \nminimum variability timescale for GRB /uniF0A0 170817A is 0.125 0.064  ( ) s.  \nUsing the soft spectral template of the GBM targeted search, a Band function ( Band et al. 1993 ) with a low energy power law index of -1.9, a high energy index of -3.7, and an E peak of 70 keV, Goldstein et al. ( 2017 ) also set 3 σ /uniFB02 ux upper limits on precursor impulsive gamma-ray emission. The limits on precursor activity out to T 200 s 0 GBM -are ( 6.8 -7.3 ) /uniF0A0 × /uniF0A0 10 -7 erg s -1 cm -2 and ( 2.0 -2.1 ) /uniF0A0 × /uniF0A0 10 -7 erg s -1 cm -2 for signals of 0.1 s and 1.0 s duration, respectively. The tail emission of GRB /uniF0A0 170817A is not consistent with the general behavior of SGRBs with extended emission ( Kaneko et al. 2015 ) . We set limits on possible extended emission over 10 s intervals out to T 100 s 0 GBM ~ + is ( 6.4 -6.6 ) /uniF0A0 × 10 -8 erg s -1 cm -2 . Additional upper limits for representative normal and harder spectra are provided in Goldstein et al. ( 2017 ) and are up to a factor of a few less constraining.", '2.3. INTEGRAL SPI-ACS Observation of GRB /uniF0A0 170817A': "The orientation of INTErnational Gamma-Ray Astrophysics Laboratory ( INTEGRAL ) with respect to the LIGO -Virgo localization of GW170817 favored the observation by SPIACS and was such that the sensitivity of the Imager on Board the INTEGRAL Satellite ( IBIS ) was much lower in comparison ( Savchenko et al. 2017b ) . For comparison of relative sensitivities of different INTEGRAL instruments see Savchenko et al. ( 2017a ) .  \nA routine follow-up search for short transients in SPI-ACS identi /uniFB01 ed a single excess at T T 1.88 0 ACS 0 GW = + s with S / N /uniF0A0 = /uniF0A0 4.6 at the 0.1 s timescale. The correction to the geocentric system assumes the location of the optical transient and results in delay of the signal arrival to INTEGRAL of 148.96 ms. In order to compare the intensity of the event observed by SPI-ACS to the GBM measurement, we compute the range of /uniFB02 uences compatible with the SPI-ACS data in the 0.320 s, 0.256 s -+ [ ] time interval centered in T 0 GBM , assuming the GBM best /uniFB01 t spectral model in the same interval. We derive a /uniFB02 uence estimate of 1.4 0.4 10 7  ' -( ) erg /uniF0A0 cm -2 ( statistical uncertainty only ) in the 75 -2000 keV energy range, consistent with GBM.  \nThe signi /uniFB01 cance of the association between the GBM observation of GRB /uniF0A0 170817A and the event observed by SPI-ACS is 4.2 σ . While SPI-ACS would not have alone reported this event as a GRB, it would have reported the event while searching around GW170817, with an independent association signi /uniFB01 cance of /uniF0A0 3.2 σ ( Savchenko et al. 2017b ) . SGRBs are routinely jointly detected by GBM and SPI-ACS and the association evidence from time coincidence ( quoted above ) as well as the consistency between the event /uniFB02 uences and temporal properties observed by the two instruments proves that both GBM and SPI-ACS observed the same event. The difference between the time of arrival of the signal in the SPI-ACS and GBM detectors can be exploited to improve the gamma-ray localization of GRB /uniF0A0 170817A, which may be bene /uniFB01 cial in future joint detections.  \nThe signi /uniFB01 cant interval of the SPI-ACS lightcurve of GRB /uniF0A0 170817A is limited to a single pulse with a duration of 100 ms ( third panel in Figure 1 ) . GBM and SPI-ACS see the main pulse as appearing to have different durations because they are sensitive in different energy ranges. If the GBM data are shown in an energy range higher than the standard 50300 keV, the main pulse is consistent with the 100 ms interval seen in SPI-ACS. The lightcurve observed by SPI-ACS reveals  \na short rise time ( 50 < ms ) and a rapid drop ( 50 < ms ) .We therefore constrain the pulse duration in the energy range observed by SPI-ACS ( ∼ 75 -2000 keV ) to less than 100 ms.", '3. Unambiguous Association': "The separation of GRBs into short and long classes was suggested by their duration distributions and reinforced by differences in the prompt gamma-ray emission of the two classes ( Dezalay et al. 1992; Kouveliotou et al. 1993 ) .Tyingtheshort class to a different progenitor from the long class was strengthened by redshift measurements of their hosts ( Berger 2014 ) . Association of SGRBs with older stellar populations than long GRBs was supported by the types of galaxies that host them ( Fong et al. 2013 ) ; the connection to BNS mergers was strengthened by the offsets of SGRBs afterglows from their host galaxies ( Trojaetal.2008;Churchetal.2011; Fong & Berger 2013 ) and by the absence of supernovae following nearby, wellobserved SGRBs ( Fox et al. 2005; Hjorth et al. 2005a, 2005b; Bloom et al. 2006; Soderberg et al. 2006;D ' Avanzo et al. 2009; Kocevski et al. 2010; Rowlinson et al. 2010;Bergeretal.2013 ) . We provide conclusive evidence for the BNS-SGRB connection by quantifying the chance temporal and spatial coincidence for GRB /uniF0A0 170817A and GW170817 arising from two independent astrophysical events.  \nTo quantify the temporal agreement, we consider the null hypothesis that SGRB and GW detection events are independent Poisson processes and determine how unlikely it is to observe an unassociated SGRB within t 1.74 s SGRB GW D = -of the GW signal. GWs from a BNS merger have been detected once to date, so the p -value is P t R 2 temporal SGRB GW GBM SGRB =D --, where R GBM SGRB -is the GBM SGRB detection rate. Using the standard duration cut T 2s 90 < , GBM triggered on-board in response to 351 SGRBs in 3324 days 184 ( the number of days between the GBM on-board trigger activation and the detection of GRB /uniF0A0 170817A ) . Further, we account for the livetime of GBM, which is disabled 15% of the time to preserve detector lifetime in regions of high particle activity during transit through the South Atlantic Anomaly. Therefore, P temporal = 2 1.74 s 351 3324 days 0.85 5.0 10 6 = ' -( )( ) , which corresponds to a 4.4 s signi /uniFB01 cance in Gaussian statistics.  \nIn order to quantify the spatial agreement of the independent GBM and LIGO -Virgo localizations, we de /uniFB01 ne the statistic PP i N i i 1 1 2 pix  = å = , where P 1 and P 2 are the posterior probabilities from GBM and LIGO -Virgo maps and i is the HEALPix ( Gorski et al. 2005 ) pixel index.  is then compared against a background distribution generated by randomly shifting and rotating GBM posteriors from a representative sample of 164 SGRBs localized by the targeted search. We factor in the estimated localization systematic, and randomly shift and rotate each map 10 times. This background method accounts for the morphology and size distributions of GBM SGRB localizations. We /uniFB01 nd a p -value P 0.01 spatial = that the two independent localizations agree this well by chance.  \nThe temporal and spatial p -values are independent quantities, thus the probability that GRB /uniF0A0 170817A and GW170817 occurred this close in time and with this level of location agreement by chance is P P temporal spatial ' = 5.0 10 0.01 6 ' ' -( ) ( ) = 5.0 10 8 ' -, corresponding to a Gaussian-equivalent signi /uniFB01 cance  \nof 5.3 s . This unambiguous association con /uniFB01 rms that BNS mergers are progenitors of ( at least some ) SGRBs.", '4. Implications for Fundamental Physics': 'Little or no arrival delay between photons and GWs over cosmological distances is expected as the intrinsic emission times are similar and the propagation speeds of EM and GWs are thought to be identical. In this Section we discuss the implications on fundamental physics of the temporal offset of 1.74 0.05 s +  ( ) measured between GW170817 and GRB /uniF0A0 170817A.  \nStandard EM theory minimally coupled to general relativity predicts that GWs and light propagate with identical speeds. The refractive index of vacuum is expected to be unity, and both waves are expected to be affected by background gravitational potentials in the same way. The arrival delay of only a few seconds across a distance greater than one hundred million light years places stringent constraints on deviations from fundamental principles. We use the observed temporal offset, the distance to the source, and the expected emissiontime difference to place constraints on the deviation of the speed of gravity from the speed of light, and on violations of Lorentz invariance and the equivalence principle.', '4.1. Speed of Gravity': "Assuming a small difference in travel time t D between photons and GWs, and the known travel distance D , the fractional speed difference during the trip can be written vv v tD EM EM D » D , where v v v GW EM D = -is the difference between the speed of gravity v GW and the speed of light v EM . This relation is less constraining for small distances, hence we conservatively use here D 26 Mpc = , the lower bound of the 90% credible interval on luminosity distance derived from the GW signal ( Abbott et al. 2017e ) . If we conservatively assume that the peak of the GW signal and the /uniFB01 rst photons were emitted simultaneously, attributing the entire 1.74 0.05 s +  ( ) lag to faster travel by the GW signal, this time difference provides an upper bound on v D . To obtain a lower bound on v D , one can assume that the two signals were emitted at times differing by more than 1.74 0.05 s +  ( ) with the faster EM signal making up some of the difference. As a conservative bound relative to the few second delays discussed in Section 2.1, we assume the SGRB signal was emitted 10 s after the GW signal. The resulting constraint on the fractional speed difference is  \nv v 3 10 7 10 . 1 15 EM 16   -' D +' --( )  \nThe intergalactic medium dispersion has negligible impact on the gamma-ray photon speed, with an expected propagation delay many orders of magnitude smaller than our errors on v GW .  \nLags much longer than 10 s are proposed in alternative models ( e.g., Ciol /uniFB01 & Siegel 2015; Rezzolla & Kumar 2015 ) , and emission of photons before the merger is also possible ( Tsang et al. 2012 ) . Hence, certain exotic scenarios can extend this time difference window to ( -100 s, 1000 s ) , yielding a 2 orders of magnitude broadening of the allowed velocity range on either side. While the emission times of the two messengers are inherently model dependent, conservative assumptions yield dramatic improvements over existing indirect ( Kostelecky  \n& Russell 2017 ) and direct ( Cornish et al. 2017 ) constraints, which allow for time differences of more than 1000 years. Future joint GW -GRB detection should allow disentangling the emission time difference from the relative propagation time, as only the latter is expected to depend on distance.", '4.2. Lorentz Invariance Violation Limits': 'Within a comprehensive effective /uniFB01 eld theory description of Lorentz violation ( Colladay & Kostelecký 1997, 1998; Kostelecký 2004; Tasson 2014 ) , the relative group velocity of GWs and EM waves, is controlled by differences in coef /uniFB01 cients for Lorentz violation in the gravitational sector and the photon sector at each mass dimension d ( Kostelecký & Mewes 2016, 2009, 2008; Wei et al. 2017 ) . We focus here on the non-birefringent, non-dispersive limit at mass dimension d /uniF0A0 = /uniF0A0 4, as it yields by far the most impressive results. In this case, the difference in group velocities for the two sectors takes the form  \nv Y n s c 1 2 1 . 2 ℓm ℓ ℓm ℓ ℓm I ℓm 2 1 4 4  å D=---+ /uni239C /uni239F /uni239B /uni239D /uni239E /uni23A0 ( ˆ) ( ) ( ) ( ) ( ) ( )  \nThe result is presented in a spherical harmonic, Yℓm ,basis, s ℓm 4 ( ) and c Iℓm 4 ( ) ( ) being spherical-basis coef /uniFB01 cients for Lorentz violation in the gravitational and EM sectors, respectively. The direction n ˆ refers to the sky position ( provided in Coulter et al. 2017a, 2017b ) .  \nFor ease of comparison with the many existing sensitivities ( Shao 2014a, 2014b; Shao et al. 2017; Kostelecký & Tasson 2015; Bourgoin et al. 2016; Le Poncin-La /uniFB01 tte et al. 2016; Kostelecky & Russell 2017 ) to the d /uniF0A0 = /uniF0A0 4 gravity-sector coef /uniFB01 cients ( Bailey & Kostelecký 2006; Hees et al. 2016 ) ,an analysis in which the coef /uniFB01 cients are constrained one at a time is useful ( Flowers et al. 2016 ) , with all other coef /uniFB01 cients, including the EM sector ones, set to zero. These results are presented in Table 1 along with the best constraints for each coef /uniFB01 cient prior to this work. These results can be compared with the isotropic A , LV a Lorentz violation parametrization ( Mirshekari et al. 2012 ) used by Abbott et al. ( 2017c ) in dispersive GW tests. The 2 LV a = limit of this parametrization is equivalent to the isotropic limit of the framework discussed above, with s A 4 00 4 p  ( ) . Constraints on A for 2 LV a = can be obtained from the /uniFB01 rst line of Table 1; these cannot be established within the analysis carried out in Abbott et al. ( 2017c ) .', '4.3. Test of the Equivalence Principle': "Probing whether EM radiation and GWs are affected by background gravitational potentials in the same way is a test of the equivalence principle ( Will 2014 ) . One way to achieve this is to use the Shapiro effect ( Shapiro 1964 ) , which predicts that the propagation time of massless particles in curved spacetime, i.e., through gravitational /uniFB01 elds, is slightly increased with respect to the /uniFB02 at spacetime case. We will consider the following simple parametrized form of the Shapiro delay ( Krauss & Tremaine 1988; Longo 1988; Gao et al. 2015; Kahya & Desai 2016 ) :  \nr t c U l dl 1 , 3 r r S 3 e o ò d g =-+ ( ( )) ( )  \nTable 1 Constraints on the Dimensionless Minimal Gravity Sector Coef /uniFB01 cients  \nNote. Constraints on the dimensionless minimal gravity sector coef /uniFB01 cients obtained in this work via Equations ( 1 ) and ( 2 ) appear in columns 3 and 5. The corresponding limits that predate this work and are reported in columns 2 and 6; all pre-existing limits are taken from Kostelecký & Tasson ( 2015 ) , with the exception of the upper limit on s 00 4 ( ) from Shao ( 2014a, 2014b ) . The isotropic upper bound in the /uniFB01 rst line shows greater than 10 orders of magnitude improvement. The gravity sector coef /uniFB01 cients are constrained one at a time, by setting all other coef /uniFB01 cients, including those from the EM sector, to zero.  \nwhere r e and r o denote emission and observation positions, respectively, r U ( ) is the gravitational potential, and the integral is computed along the wave path. γ parametrizes a deviation from the Einstein -Maxwell theory, which minimally couples classical electromagnetism to general relativity. We allow for different values of γ for the propagation of EM and GWs ( EM g and GW g , respectively, with 1 EM GW g g = = in the Einstein -Maxwell theory ) .  \nWhile obtaining the best bound on the difference between the Shapiro time delays requires modeling the potential r U ( ) along the entire line of sight, we determine a conservative bound on GW EM g g -by considering only the effect of the Milky Way outside a sphere of 100 kpc, and by using a Keplerian potential with a mass of M 2.5 10 11 '  ( the lowest total mass within a sphere of radius 100 kpc quoted in BlandHawthorn & Gerhard 2016, from Gibbons et al. 2014, taking the 95% con /uniFB01 dence lower bound )( Krauss & Tremaine 1988; Longo 1988; Gao et al. 2015 ) . Using the same time bounds as Equation ( 1 ) we /uniFB01 nd  \n2.6 10 1.2 10 . 4 7 GW EM 6   g g -' -' --( )  \nThe best absolute bound on EM g is 1 2.1 2.3 EM g - =  ' ( ) 10 5 -, from the measurement of the Shapiro delay ( at radio wavelengths ) with the Cassini spacecraft ( Bertotti et al. 2003 ) .", '5. Astrophysical Implications': 'The joint GW -GRB detection provides us with unprecedented information about the central engine of SGRBs. The delay between the GW and the GRB trigger times allows us to examine some basic GRB physics. This delay could be intrinsic to the central engine, re /uniFB02 ecting the time elapsed from the moment the binary components come into contact to the formation of a remnant BH and the resulting jet. This interpretation includes the case of a relatively long-lived massive NS remnant, which has been suggested to survive from seconds to minutes after merger /uniF0A0 ( see Faber & Rasio 2012; Baiotti & Rezzolla 2017 and references therein ) . The delay could also be due to the propagation time of the relativistic jet,  \nincluding the time it takes for the jet to break out of the dense gaseous environment produced by non-relativistic merger ejecta /uniF0A0 ( Nagakura et al. 2014; Moharana & Piran 2017 ) and / or the emitting region to become transparent to gammarays /uniF0A0 ( Mészáros & Rees 2000 ) .  \nWe /uniFB01 rst discuss the implications that the time delay between the GW and EM emission has on the physical properties of the emitting region when considering the jet propagation and transparency scenarios. Here we assume that the entire delay is due to the expansion of the emitting region and neglect any intrinsic delays between the moment of binary coalescence and the launching of the resulting jet, thus placing limits on the physical properties of the system. Then we consider the impact of SGRB emission from an NS merger on the EOS of dense matter.', '5.1. GRB Physics': "The main hard peak observed for GRB /uniF0A0 170817A lasted roughly half a second. This peak is consistent with a single intrinsic emission episode as it is well described by a single pulse ( Goldstein et al. 2017 ) , showing no evidence for signi /uniFB01 cant substructure ( spikes ) . This interpretation is consistent with the SPI-ACS observation of a single peak ( Savchenko et al. 2017b ) . The GBM detection of GRB /uniF0A0 170817A also shows no evidence for photons with energy > 511 keV, implying that the out /uniFB02 ow does not require a high bulk Lorentz factor Γ to overcome photon -photon absorption at the source.  \nExplanations for the extreme energetics and short timescales observed in GRBs invoke a near instantaneous release of a large amount of energy in a compact volume of space /uniF0A0 ( Goodman 1986; Paczynski 1986 ) . This is commonly referred to as the /uniFB01 reball model, and it is the framework that we will assume for the remainder of this section. The /uniFB01 reball model is largely independent of the burst progenitor and focuses on the dynamics of such a system after this sudden release of energy. The resulting pair-plasma is optically thick and quickly expands under its own pressure to produce a highly relativistic out /uniFB02 ow that coasts asymptotically with a constant Lorentz factor Γ .Withinthe /uniFB01 reball, kinetic energy is imparted to particles  \nentrained in the out /uniFB02 ow, although alternative models exist in which the energy out /uniFB02 ow occurs mostly as Poynting /uniFB02 ux ( Usov 1992; Lyutikov & Blandford 2003 ) . The observed gamma-ray pulses are attributed to shocks internal to this relativistic out /uniFB02 ow, which convert some of their kinetic energy into the observed EM radiation /uniF0A0 ( Rees & Meszaros 1994 ) . These shocks could produce the predominantly non-thermal emission observed in most GRBs, although non-shock heating models have also been proposed /uniF0A0 ( e.g., Giannios 2006 ) . The overall multi-pulse duration of a burst is thought to re /uniFB02 ect the time that the inner engine was active ( e.g., producing inhomogeneities in the out /uniFB02 ow represented as shells traveling with different bulk velocities ) and the variability of individual pulses re /uniFB02 ects the size of the shells producing the emission. For a top-hat jet model, Γ is assumed constant over the jet surface and the observer never sees beyond the beaming angle 1 b q ~ G . Therefore, the values inferred from the data are independent of the inclination angle from the total angular momentum axis of the system, as long as the viewer is within the opening angle of the jet.  \nWe can examine the implications of the observed delay between the GW and EM signals in the internal shock scenario if we consider two shells emitted at time t 0 GW = and time t GW + t engine D . If the Lorentz factor of the second shell, 2 g ,is greater than the Lorentz factor of the /uniFB01 rst shell, 1 g , the shells will collide at time  \nt t 1 , 5 delay engine 1 2 2 gg = D -( ) ( )  \nwhich is valid if , 1 1 2 g g  . If the shells have comparable masses, conservation of energy and momentum leads to a merged shell with Lorentz factor m 1 2 1 2 g gg = ( ) . The resulting pulse pro /uniFB01 le is determined by two timescales. The rise time ( which we equate to the minimum variability timescale ) can be attributed to the light-crossing time of the individual emission regions and is expressed as  \nt R c 2 , 6 rise m 2 d g D » ( )  \nwhere R d is the thickness of the emitting region. The decay time re /uniFB02 ects angular effects, where off-axis emission is delayed and affected by a varying Doppler boost due to the curvature of the relativistic shell. This timescale is essentially the difference in light-travel time between photons emitted along the line of sight and photons emitted at an angle θ along a shell of radius R . This timescale may be expressed as  \nt R c R c R c t 1 cos 2 2 , 7 decay 2 m 2 rise q q g D = -D » D » > D ( ) ( ) ( )  \nwhere we assume that the solid angle accessible to the observer is limited by relativistic beaming and thus given by θ /uniF0A0 ∼ /uniF0A0 1 / γ . At the same time, the distance that the /uniFB01 rst shell has traveled since ejection is R c t 2 1 1 2 delay g » , leading to  \nt t t 1 . 8 decay delay 1 2 engine 1 2 2 1 2 g g gg g g D » = D -( ) ( ) ( )  \nThe conclusion is a linear correlation between the delay in the GW and EM signals and the resulting pulse duration, modulo  \nthe ratio of the Lorentz factors of the two colliding shells /uniF0A0 ( Fenimore et al. 1996; Kocevski et al. 2007; Krimm et al. 2007 ) .  \nThe relative similarity between the gamma-ray duration T 90 and the delay between the GW and the EM emission gives t t 1 decay delay D ~ , pointing to an internal shock scenario in which the difference in the Lorentz factors of the colliding shells, g D , is much smaller than their typical values, i.e., g g D  . This would imply that the jet was launched shortly after the time of the merger and points to a relatively short t engine D time in which the central engine was active. Such a scenario would produce a collision that was relatively inef /uniFB01 cient at converting the internal energy of the shocks to radiation, resulting in a signi /uniFB01 cant isotropic equivalent kinetic energy remaining in the merged shell /uniF0A0 ( Kobayashi et al. 1997; Krimm et al. 2007 ) . This would lead to a very signi /uniFB01 cant energy injection into the resulting afterglow, producing late time ' refreshed shocks ' ( Rees & Mészáros 1998; Kumar & Piran 2000; Ramirez-Ruiz et al. 2001 ) , which are typically not observed in the X-ray ( and optical ) lightcurves of SGRBs ( e.g., Perley et al. 2009 ) .  \nSome of these energetics constraints can be alleviated if we exclude the soft thermal emission from the gamma-ray duration estimate. In this case, the prompt non-thermal emission of t 0.5 s decay D  would be due to internal shocks and the soft thermal emission would be attributed to a separate component. In this case we obtain t t 0.3 decay delay D  , implying 3 2 1 g g » . These energetics considerations may suggest that the initial hard pulse and the subsequent thermal emission observed by GBM mayindeedbedistinctcomponents.  \nWithin the context of the internal shock model, if we assume the entire 1.74 0.05 s +  ( ) delay between the GW and the EM emission is due to jet propagation time and use a Lorentz factor of 100 g < for the /uniFB01 rst shell, we can estimate an upper limit to the radius of the relativistic out /uniFB02 ow to be R 5 10 14 ~ ' cm or ∼ 30 au. The minimum variability timescale t t 0.125 s rise min D =D ~ /uniF0A0 ( Goldstein et al. 2017 ) yields an upper limit on the size of the emitting region of R 4 10 13 d ~ ' cm, or ∼ 3 au. The ratio of the two is independent of the unknown Lorentz factor and is of order RR 10% d ~ .  \nThe single-pulsed nature of the gamma-ray emission, as well as the observed t t 1 decay delay D ~ , also leaves open the possibility that the GBM signal is entirely of an external shock origin. In this scenario, the relativistic out /uniFB02 ow converts its internal energy to radiation due to its interaction with an external medium, such as the interstellar matter /uniF0A0 ( Meszaros & Rees 1992 ) . If we associate the duration of the main pulse with the deceleration time, i.e., the timescale over which the jet is signi /uniFB01 cantly decelerated by interstellar matter of constant density n , in the external shock scenario ( Dermer et al. 1999 ) :  \nt E nm c T 3 4 . 9 dec k,iso 8 p 5 1 3 obs pg = = [ ( )] ( )  \nE k,iso is the kinetic energy of the jet calculated assuming a gamma-ray production ef /uniFB01 ciency of 20%, m 1.67 10 p 27 = ' -kg is the proton mass, c is the speed of light, and T obs is the approximate duration of the main peak. We can thus estimate  \nthe Lorentz factor of the jet in the external shock scenario to be  \nE n T 310 2 10 erg 0.1 cm 0.5 s . 10 k,iso 47 1 8 3 1 8 obs 3 8 g » ' ---/uni239C /uni239F /uni239C /uni239F /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F /uni239B /uni239D /uni239E /uni23A0 /uni239B /uni239D /uni239E /uni23A0 ( )  \nThe deceleration radius represents the upper limit of ef /uniFB01 cient energy extraction ( even for internal shocks ) and can be expressed as  \nR cT E n T 2 3.0 10 2 10 erg 0.1 cm 0.5 s cm. 11 dec 2 obs 15 k,iso 47 1 4 3 1 4 obs 1 4 g = = ' ' ' --/uni239C /uni239F /uni239C /uni239F /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F /uni239B /uni239D /uni239E /uni23A0 /uni239B /uni239D /uni239E /uni23A0 ( )  \nTherefore, the deceleration radius and associated Lorentz factor also serve as upper limits to the radius and Lorentz factor of the emitting region in the internal shock scenario.  \nThe soft thermal component observed by GBM could also be due to the photosphere of the /uniFB01 reball before it becomes optically thin to gamma-rays. In this interpretation, the delay between the GW and the GRB trigger times may represent the time it takes for the relativistic /uniFB01 reball to expand and become optically thin to gamma-ray radiation. We can examine this scenario by estimating the time it takes for a /uniFB01 reball to become transparent to high-energy radiation in an environment similar to that of a BNS merger.  \nFollowing Mészáros & Rees ( 2000, hereafter MR00 ) , we assume an out /uniFB02 ow with an initial radius R GM c 6 0 BH 2 = ( the innermost stable circular orbit of a Schwarzschild BH with a mass equal to M BH ) .Inourcase R 2.5 10 0 6 = ' cm, and M BH = M 2.8  . Given the GBM observations, we estimate an isotropic equivalent energy of the soft thermal BB component to be E 1.3 10 iso,BB 46 = ' erg and peak isotropic luminosity of L iso,p = 1.6 10 47 ' erg s -1 ( see Section 6.1 ) . We take this luminosity as an upper bound of the average luminosity, which may be estimated as L E t 1.1 10 iso,BB iso,BB BB 46 = D = ' erg s -1 , where we have used a duration of the soft BB component of t 1.15 s BB D = .  \nUsing these parameters along with /uniFB01 ducial values ( see Appendix B for details ) , we estimate the photosphere radius to be ( MR00 )  \nR L Y mc L Y 4 2.01 10 cm 10 erg s 18 . 12 T ph p 3 3 13 0 50 1 3 s p h h = = ' ' --/uni239C /uni239F /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F /uni239B /uni239D /uni239E /uni23A0 ( )  \nWhere L 0 is the initial /uniFB01 reball luminosity, Y is the number of electrons per baryon ( in our case Y 1  ) , 6.65 10 T 25 s = ' -cm 2 is the Thomson cross-section, and η is the dimensionless entropy of the /uniFB01 reball, whose value is much smaller than the canonical one in the standard /uniFB01 reball model ( see Appendix B ) . We note that L 0 can be much larger than L iso,BB , since the /uniFB01 reball must expand and convert the remaining internal energy into kinetic energy of the ejecta.  \nThe laboratory frame time needed for the /uniFB01 reball expanding at roughly the speed of light to reach the transparency radius is  \nt R c ph ph  ; thus,  \nt L Y 672 s 10 erg s 18 . 13 ph 0 50 1 3 h -- /uni239C /uni239F /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F /uni239B /uni239D /uni239E /uni23A0 ( )  \nFollowing Bianco et al. ( 2001 ) , we can set a upper bound to the conversion from the laboratory to observer frame by assuming the observer is viewing the /uniFB01 reball at most at an angle v c cos J = :  \nt t L Y 2.1 s 10 erg s 18 . 14 a ph 2 0 50 1 5 g h --  /uni239C /uni239F /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F /uni239B /uni239D /uni239E /uni23A0 ( )  \nThis upper bound can account for the time delay between the GW and the prompt radiation in the soft thermal peak.  \nEmploying MR00 ' s Equation ( 8 ) we can estimate the observer frame temperature of the expanding /uniFB01 reball at the photospheric radius  \nT T L R Y 2.3 keV 10 erg s 2.5 10 cm 18 . 15 ph obs ph 0 50 1 5 12 0 6 5 6 11 3 2 3 0 5 6 h h g = ' ' ----- /uni239C /uni239F /uni239C /uni239F /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F /uni239B /uni239D /uni239E /uni23A0 /uni239B /uni239D /uni239E /uni23A0 ( )  \nT ph obs can then be compared to the one obtained from observational /uniFB01 ts to the GBM data, which provide a BB temperature of T 10.3 BB obs  keV. Our result underestimates the observed BB temperature by a factor ∼ 4, but we are neglecting Comptonization effects, which may slightly raise the estimated temperature. The corresponding BB luminosity is ( MR00 )  \nL L R Y 1.7 10 erg s 10 erg s 2.5 10 cm 18 . 16 ph 46 1 0 50 1 7 12 0 6 5 6 8 3 2 3 0 5 6 h g ' ' ' ----- /uni239C /uni239F /uni239C /uni239F /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F /uni239B /uni239D /uni239E /uni23A0 /uni239B /uni239D /uni239E /uni23A0 ( )  \nAs we have mentioned above, the average luminosity of the BB component is L 1.1 10 iso,BB 46 '  erg s -1 , which is of the order of L ph estimated here.  \nTherefore, based on the observed temperature and luminosity, the delay between the GW signal and the soft BB component can be accounted for as the time it takes the /uniFB01 reball photosphere to radiate. The primary challenge of this interpretation is in explaining the nature of the hard nonthermal emission preceding the BB component. If both components are the result of the same expanding /uniFB01 reball, the photospheric emission is expected to occur earlier than or at the same time as the non-thermal emission. This requirement can be reconciled with the GBM data if the thermal component was subdominant and indistinguishable during the initial hard nonthermal pulse.  \nAlternatively, energy dissipation below the photospheric radius could also provide an explanation for the timing of the two pulses /uniF0A0 ( Rees & Mészáros 2005 ) . This could be achieved through a range of possible scenarios. Energy dissipation could occur through inelastic collisions between the decoupled neutron and proton populations within the jet ( Beloborodov 2010 ) , for example, or through magnetic reconnection processes ( Giannios 2006 ) . The emitted radiation would exhibit a modi /uniFB01 ed blackbody spectrum and be released at the  \n<!-- image -->  \n<!-- image -->  \nFigure 3. Critical mass boundaries for different EOSs in comparison with the 90% credible region of the gravitational masses inferred from GW170817 ( prior limits on the spin magnitude, z c /uni2223 /uni2223 , given in the legend ) . The slanted curves in the left panel and middle panel correspond to the maximum baryonic mass allowed for a single non-rotating NS ( left ) and for a uniformly rotating NS ( middle ) . Arrows indicate for each EOS the region in the parameter space where the total initial baryonic mass exceeds the maximum mass for a single non-rotating or uniformly rotating NS, respectively. The right panel illustrates EOS-dependent cuts on the gravitational mass m 1 of the heavier star, with arrows indicating regions in which m 1 exceeds the maximum possible gravitational mass M G Static for non-rotating NSs. In all three panels the black solid line marks the m m 1 2 = boundary, and we work in the m m 1 2 > convention.  \n<!-- image -->  \nphotospheric radius, resulting in similar arrival times to the thermal emission. A non-thermal pulse could also arise from a forward shock if the deceleration radius were located below the photosphere. Such a scenario would be possible if the density of the external medium were suf /uniFB01 ciently in excess of the interstellar medium, which is a distinct possibility for such environments /uniF0A0 ( Goriely et al. 2011; Bauswein et al. 2013; Hotokezaka et al. 2013 ) .  \nThe thermal component could also be the result of ' cocoon ' emission from shocked material surrounding the relativistic jet /uniF0A0 ( Lazzati et al. 2017 ) , which is expected to be softer and fainter than the non-thermal prompt emission ( Ramirez-Ruiz et al. 2002;Pe ' er et al. 2006 ) . To examine this scenario, we utilize the relation between the radius of the shock breakout, duration and observed temperature proposed by Nakar & Sari ( 2012 ) :  \nR t T 1.4 10 1 s 10 keV cm. 17 9 2 = ' /uni239C /uni239F /uni239C /uni239F /uni239B /uni239D /uni239E /uni23A0 /uni239B /uni239D /uni239E /uni23A0 ( )  \nUsing the distance measurement, the blackbody spectral /uniFB01 t implies a radius of R 3 10 cm BB 8 = ' that we can use as a proxy for the cocoon radius. This is within a factor of 4 of the relation, in spite of the fact that it applies to spherical geometry that is not guaranteed here, and the R BB radius derived from the /uniFB01 t assumes thermal equilibrium. We thus consider this as evidence supporting the cocoon scenario.  \nFinally, the delay between the GW and EM signals may also be due in part to the time it takes for the relativistic jet to break out of the sub-relativistic dense ejecta surrounding the merger /uniF0A0 ( Nagakura et al. 2014; Moharana & Piran 2017 ) .We estimate that the breakout time for typical dynamical ejecta mass values of M 0.1 ~  in such a merger /uniF0A0 ( Hotokezaka et al. 2013 ) could not account for the entire observed delay. Lowering the reference isotropic kinetic luminosity of L 10 k ,iso 51 = erg s -1 assumed by Moharana & Piran ( 2017 ) could be one way to account for a larger delay. However, a luminosity below the one assumed in the breakout scenario substantially increases the likelihood of a ' choked ' jet that fails to break out of the surrounding medium /uniF0A0 ( Aloy et al. 2005 ) .", '5.2. Neutron Star EOS Constraints': 'The observation of an SGRB associated with the merger of two NSs can be used to derive constraints on the EOS of NS matter ( see theoretical studies by Belczynski et al. 2008; Fryer et al. 2015; Lawrence et al. 2015 ) . To do this, we compare the measurement of the binary mass from the GW signal with two possible models of the merger remnant that powered the SGRB: ( i ) the merger remnant collapsed to a rotating BH with a surrounding disk that powered the SGRB /uniF0A0 ( Shibata et al. 2006 ) ,or ( ii ) the merger formed a rapidly rotating, strongly magnetized NS ( millisecond magnetar ) with an accretion disk ( Metzger et al. 2008 ) .  \nWe consider a representative sample of EOSs: SLy ( Douchin & Haensel 2001 ) , LS220 ( Lattimer & Swesty 1991 ) , SFHo ( Steiner et al. 2013 ) ,H4 ( Lackey et al. 2006 ) , APR4 ( Akmal et al. 1998 ) , SHT ( Shen et al. 2011 ) , and MS1 ( Müller & Serot 1996 ) . For each EOS, we compute the maximum stable baryonic mass and gravitational mass of a non-rotating ( static ) NS, denoted M B Static and M G Static , respectively, and the maximum baryonic mass of a uniformly rotating NS M B Uniform ( Gourgoulhon et al. 2001 ) . The merger remnant can only collapse to a BH if its baryonic mass is larger than M B Static .  \nIf we neglect rotational corrections, the baryonic masses m B1 , m B2 of the initial NSs are functions of their gravitational masses m 1, m 2 only. In this approximation, a /uniFB01 xed total initial baryonic mass, M B Initial , corresponds to a curve in the ( m m , 1 2 ) parameter space. In Figure 3 we show lines of M B Static and M B Uniform that bound the region of the parameter space in which the total mass of the binary is consistent with a stable non-rotating or uniformly rotating remnant, respectively. The /uniFB01 gure also contains the 90% credible region of the gravitational masses obtained with a restricted or full spin prior ( Abbott et al. 2017e ) . We note that the latter has a broader distribution of the component masses, such that the heavier NS can exceed M G Static for various EOS, which would correspond to either a supramassive ( or even hypermassive ) NS, or to a light BH. The maximum gravitational masses allowed for each EOS, M G Static , are shown in the /uniFB01 gure as vertical lines.  \n<!-- image -->  \nFigure 4. GRB /uniF0A0 170817A is a dim outlier in the distributions of E iso and L iso , shown as a function of redshift for all GBM-detected GRBs with measured redshifts. Redshifts are taken from GRBOX ( http: // www.astro.caltech.edu / grbox / grbox.php ) and Fong et al. ( 2015 ) . Short- and long-duration GRBs are separated by the standard T 2s 90 = threshold. For GRBs with spectra best modeled by a power law, we take this value as an upper limit, marking them with downward pointing arrows. The power law spectra lack a constraint on the curvature, which must exist, and therefore, will overestimate the total value in the extrapolated energy range. The green curve demonstrates how the ( approximate ) GBM detection threshold varies as a function of redshift. All quantities are calculated in the standard 1 keV -10 MeV energy band.  \n<!-- image -->  \nSince the total baryonic mass of the system can only be reduced ( by mass ejection ) , the maximum baryonic mass of the merger remnant and accretion disc is bound by M B Initial . From Figure 3, we can see that for the measured NS gravitational masses with the low-spin prior, the MS1 and SHT EOS could not form a BH since M M B Initial B Static < . Assuming that the magnitude of the spins is small, the MS1 and SHT EOS are incompatible with BH formation. If the dimensionless spins of the NSs are allowed to be larger than 0.05, BH formation is only disfavored: we /uniFB01 nd that a fraction 83% ( MS1 ) and 84% ( SHT ) of the posterior distribution satis /uniFB01 es M M B Initial B Static < . For both spin priors, we /uniFB01 nd that the H4, LS220, SFHo, and SLy EOS result in M M B Initial B Uniform > . Even when assuming a large ejecta mass of M 0.1  , the remaining mass cannot form a uniformly rotating NS. For those EOS, the merger either results in prompt BH formation or in a short-lived remnant, with a lifetime determined by the dissipation of differential rotation and / or disk accretion.  \nSection 6.5 we will discuss some model-dependent implications of the lack of precursor and temporally extended gamma-ray emission from GRB /uniF0A0 170817A on the progenitor NSs.', '6. Gamma-ray Energetics of GRB /uniF0A0 170817A and their Implications': 'Using the measured gamma-ray energy spectrum and the distance to the host galaxy identi /uniFB01 ed by the associated optical transient, we compare the energetics of GRB /uniF0A0 170817A to those of other SGRBs at known redshifts. Finding GRB /uniF0A0 170817A to be subluminous, we discuss whether this dimness is an expected observational bias for joint GW -GRB detections, what insight it provides regarding the geometry of the gammaray emitting region, what we can learn about the population of SGRBs, update our joint detection estimates, and set limits on gamma-ray precursor and extended emission.', '6.1. Isotropic Luminosity and Energetics of GRB /uniF0A0 170817A': "Using the ' standard ' spectral information from Goldstein et al. ( 2017 ) and the distance to the host galaxy NGC 4993 42.9 3.2  ( ) Mpc, we calculate the energetics of GRB /uniF0A0 170817A using the standard formalisms ( Bloom et al. 2001; Schaefer 2007 ) . GRBs are believed to be relativistically beamed and their emission collimated ( Rhoads 1999 ) . Isotropic energetics are upper bounds on the true total energetics assuming the GRB is observed within the beaming angle of the brightest part of the jet. We estimate that the isotropic energy release in gamma-rays E 3.1 0.7 10 iso 46 =  ' ( ) erg, and the isotropic peak luminosity, L 1.6 0.6 10 iso 47 =  ' ( ) erg s -1 , in the 1 keV -10 MeV energy band. These energetics are from the source interval -i.e., the selected time range the analysis is run over -determined in the standard manner for GBM spectral catalog results, allowing us to compare GRB /uniF0A0 170817A to other GRBs throughout this section. The uncertainties on the inferred isotropic energetics values here include the uncertainty on the distance to the host galaxy. As a cross check, the isotropic luminosity is also  \nTo be compatible with scenario ( ii ) , the lifetime of the merger remnant would have to be suf /uniFB01 ciently long to power the GRB. We note that prompt BH formation is a dynamic process accessible only to numerical relativity simulations. Although there are parameter studies ( Hotokezaka et al. 2011; Bauswein et al. 2013 ) , they only consider equal mass binaries. Considering also the error margins of those studies, we currently cannot exclude prompt collapse for the H4, LS220, SFHo, and SLy EOS. Finally, we note that for the APR4 EOS only the possibility of a stable remnant can be ruled out. More generally, only EOSs with M M 3.2 B Static <  are consistent with scenario ( i ) when assuming the low-spin prior, or with M M 3.7 B Static <  for the wider spin prior. These bounds were derived from the 90% credible intervals of the M B Initial posteriors ( and these, in turn, are determined for each EOS in order to account for binding energy variations ) . These upper limits are compatible with and complement the lower bounds on M G Static from the observation of the most massive known pulsar, which has a mass of M 2.01 0.04   ( ) ( Antoniadis et al. 2013 ) .In  \n<!-- image -->  \n<!-- image -->  \nFigure 5. Three potential jet viewing geometries and jet pro /uniFB01 les that could explain the observed properties of GRB /uniF0A0 170817A, as described by scenarios ( i ) -( iii ) in Section 6.2.  \n<!-- image -->  \nestimated using a Bayesian approach proposed by Fan ( 2017 ) . Assuming a /uniFB02 at prior on isotropic luminosity, we obtain L iso = 1.2 10 0.6 0.7 47 ' -+ erg s -1 , which is consistent with the standard GBMapproach. This Bayesian approach can be used to combine future joint GW-GRB observations to provide a redshiftindependent estimate of the GRB luminosity function.  \nThe two apparent components of GRB /uniF0A0 170817A are suf /uniFB01 ciently different that using an average spectrum to estimate the /uniFB02 uence may produce an inaccurate total luminosity. Therefore, we also estimate E iso using the ' detailed ' /uniFB01 ts described in Goldstein et al. ( 2017 ) . Separating the hard peak best /uniFB01 t by a Comptonized function ( a power law with an exponential cutoff ) and the softer tail best /uniFB01 t by a BB spectrum, we estimate E 4.0 1.0 10 iso,comp 46 =  ' ( ) /uniF0A0 erg, and E iso,BB = 1.3 0.3 10 46  ' ( ) /uniF0A0 erg, for a total of E 5.3 iso =  ( 1.0 10 46 ' ) /uniF0A0 erg.  \nCompared to the distribution of GBM detected GRBs with measured redshift shown in Figure 4,GRB /uniF0A0 170817A is 2 orders of magnitude closer and 2 to 6 orders of magnitude less energetic than other SGRBs. In particular, GRB 150101B was previously the weakest SGRB with a /uniFB01 rm redshift association ( z 0.134; = Fong et al. 2016 ) , and its energetics ( as measured by GBM ) E 2.3 10 iso 49 = ' erg, and L 7.5 10 iso 49 = ' erg s -1 are 2 -3 orders of magnitude higher. As this was the previous dimmest burst, the minimum luminosity cut of 5 10 49 ' erg s -1 used in Wanderman & Piran ( 2015 ) to /uniFB01 tarateandan L iso distribution to existing observations appeared reasonable; however, with GRB /uniF0A0 170817A, the lower bound on the isotropic energetics distributions needs to be revised, as discussed in Section 6.4.", '6.2. Implications of the Dimness on the Central Engine': "The broad observed brightness distribution likely arises from a mixture of an intrinsic brightness distribution and geometric effects, which include the inclination angle of the system to Earth, the structure and width of the collimated jet itself, and the relativistic beaming angle b q . We consider several possibilities to explain why GRB /uniF0A0 170817A is extremely dim ( Figure 5 ) : ( i ) we viewed it from beyond the half-jet opening angle j q for a standard top-hat model, ( ii ) the structure of the jet is more complicated than a simple top-hat model, ( iii ) the observed emission for GRB /uniF0A0 170817A originates from a  \ndifferent mechanism than for most SGRBs, or ( iv ) it is due solely to the intrinsic luminosity distribution and not the geometry of the system.  \nScenario ( i ) . Uniform top-hat jets ( constant emissivity and Lorentz factor, Γ , within the jet aperture ) with a sharp edge have been widely used to explain GRB properties, including jet breaks ( Rhoads 1999 ) . The top-hat jet is the simplest possible model for calculating off-axis parameters as it captures the basic physics of the system, but it is unable to account for smooth pro /uniFB01 les in the Lorentz factor and the emissivity. Here the observed energetics are signi /uniFB01 cantly lower than they would be if we were within j q .  \nIn the top-hat scenario, off-axis values of physical quantities can be related to the on-axis values through the angle dependence of the relativistic Doppler factor:  \n1 cos 2 1 , 18 D 1 2 2 dq b q q =G »G + G -( ) [ ( )] ( ) ( )  \nwhere θ is the angle between the velocity vector v and the line of sight, and vc b = . The relation for duration and peak energy is linear with D d ( see, e.g., Granot et al. 2002 ) :  \n19 T T E E b off axis on axis on axis off axis 0 1 cos 1 1 , 90 90 p p D D j j 2 j 2 d dq z b q z b z q --= --= -= ---= » + G -D ( ) ( ) ( ) ( ) ( ) ( ) ( ) ( ) ( )  \nwhereas E off axis ,iso g ( /uni2010 ) scales approximately b 2 µ -for a viewing angle ζ between j q and 2 j q . The duration in the onaxis scenario may be longer than inferred from the above equation, as the variable gamma-ray /uniFB02 ux can be discerned above detector noise for a longer fraction of the total activity compared to emission viewed off-axis.  \nWe use the observed quantities for GRB /uniF0A0 170817A, E 200 keV p » , E 5.3 10 ,iso 46 = ' g erg, and T 2s 90 » , as values observed off-axis. If we assume that the on-axis values for GRB /uniF0A0 170817A are consistent with typical values observed for SGRBs, we obtain E b 6 30 p = ( ) MeV, E b 5 10 30 ,iso 49 2 = ' g ( ) erg, and T b 7 10 30 90 2 1 = ' --( ) s. In particular using a /uniFB01 ducial range on E on axis ,iso -g ( ) corresponding to the two orders of magnitude spread shown in  \nFigure 4 we obtain b 30 2 j 2 z q »G -» ( ) within a factor 3, which is a constraint on the values of Γ , ζ and j q .  \nIf we assume a viewing angle of 30 z =  and 300 G = the uncertainty on b yields 1 0.5 j z q -  deg, a solid angle covering only 1% of a full sphere. Hence this con /uniFB01 guration would require a /uniFB01 ne tuning of the line of sight. However, if we assume 30 G = then the uncertainty on b yields 10 4 j z q -  deg, a solid angle that covers 10% of a full sphere, which is plausible without too much /uniFB01 ne tuning. This argument only weakly depends on the particular value ζ , and illustrates that for large Γ a top-hat jet scenario is disfavored due to the sharp emission fall-off at the edges.  \nScenario ( ii ) . A more complex geometry involves a structured jet ( Rossi et al. 2002, or Granot 2007 and references therein ) which provides a wider range of angles from which the observer could still detect emission, and therefore does not require a /uniFB01 ne-tuned viewing angle. Structured jet emission pro /uniFB01 les include a uniform ultra-relativistic core surrounded by a power-law decaying wing where the energy and Lorentz factor depend on the distance from the jet axis ( Pescalli et al. 2015 ) , a Gaussian with a smooth edge and falloff outside the core ( Zhang & Mészáros 2002; Kumar & Granot 2003 ) , and a two-component jet with an ultra-relativistic narrow core and slightly slower outer jet ( Frail et al. 2000; Berger et al. 2003; Racusin et al. 2008; Filgas et al. 2011 ) , among other possibilities.  \nStructured jets can naturally explain the broad observed energetics distribution. Because SGRBs involve relativistic velocities, radiation is strongly beamed into angle b q . If the observed brightness depends on viewing geometry, i.e., is not uniform across the angle j q , then the part of the beam that we observe may be off-axis to the brightest part of the jet but we may still be within b q of some dimmer part of the emitting region, though in this case we would expect the Γ factor to vary as well.  \nScenario ( iii ) . Given the closeness of this burst it is possible that the observed emission is due to a different mechanism from other SGRBs, one that is intrinsically dim and thus undetectable at usual SGRB distances. We believe this explanation to be unlikely as the main emission episode of GRB /uniF0A0 170817A is a typical SGRB ( as measured by the observed gamma-ray properties ) . It is possible that the soft tail emission arises from a distinct mechanism. One explanation is ' cocoon ' emission from the relativistic jet shocking its surrounding non-relativistic material ( Lazzati et al. 2017 ) .We showed that ' cocoon ' emission could explain the thermal tail in Section 5.1. A possible full model for GRB /uniF0A0 170817A is offaxis emission from a top-hat jet providing the main emission episode, with ' cocoon ' emission arising from the jet ' s interaction with the surrounding torus that powers the main jet. The softer emission is near the detection limits of GBM and would not be detected to much greater distances, suggesting it may be a common property of SGRBs that is otherwise missed.  \nScenario ( iv ) . If GRB /uniF0A0 170817A is viewed within both the collimated jet and the beaming angle, and the emission is constant across the traditional top-hat jet, then GRB /uniF0A0 170817A is intrinsically much dimmer by orders of magnitude compared to other observed GRBs. This would mean that top-hat jets have an intrinsic distribution covering 6 orders of magnitude, which is dif /uniFB01 cult to envision given the limited mass ranges in the merger of two NSs ( although see Metzger & Berger ( 2012 )  \nand references therein ) . A broader intrinsic luminosity distribution might be accommodated if we assume that at least some SGRBs arise from the merging of an NS with a BH. It is possible, for example, that the brightest events may arise from NS -BH mergers with optimal mass ratio and spin parameters. Another possibility is that this broad luminosity range could arise from other properties of the system, such as the magnetic /uniFB01 eld strength of the progenitors or the intrinsic jet-opening angle distribution.  \nObservations of GW170817 / GRB /uniF0A0 170817A at other wavelengths ( which are not explored in this Letter ) will be necessary for a full understanding of this event. For example, evidence for X-ray emission that only arises at late times may provide evidence for this event occurring off-axis ( see, e.g., Mészáros et al. 1998; Granot et al. 2002; Yamazaki et al. 2002 ) . However, future joint detections of GW-GRB events can also provide a fuller understanding of the intrinsic energetics distributions and the effect geometry has on our observed brightness. Here the inclination constraint is not particularly informative as the inclination angle constraint, 36  z  , is comparable to the highest lower limit for a half-jet opening angle, 25 j q > deg ( Fong et al. 2015 ) . If this is truly off-axis from a top-hat jet then it is unlikely to be a common occurrence. Only joint GW -EM detections will reveal if the intrinsic brightness varies according to the type of progenitor.  \nThe updated expected joint detection rates in Section 6.4 suggest inferences on populations of joint detections may be possible sooner than previously thought.", '6.3. Observational Bias Against Low-luminosity GRBs': "The fact that GRB /uniF0A0 170817A is orders of magnitude dimmer than the population of SGRBs with known redshifts raises the questions: ( i ) is it unexpectedly dim, and ( ii ) is there a population of SGRBs with comparable luminosities ( and distances ) that we are not detecting? We explore here whether the gap in luminosity compared to more luminous SGRBs is a result of the instrumental sensitivity for the detection of either the prompt or the afterglow emission of SGRBs, or whether our problem lies in the association of SGRBs to their host galaxies and thus redshift.  \nBurns et al. ( 2016 ) examine the observed relationship between redshift and gamma-ray /uniFB02 uence for SGRBs with known redshift and /uniFB01 nd no strong correlation. SGRBs that appear extremely bright are likely to be nearby because their inferred luminosities would otherwise be unrealistic, but SGRBs near the detection threshold of ( current ) GRB detectors are as likely to be nearby as far away. The intrinsically dim part of the SGRB luminosity distribution is detectable only at short distances.  \nGRB /uniF0A0 170817A is our only clear case of a subluminous SGRB with known distance, so we investigate the maximum distance at which it could have triggered GBM. Assuming the event occurred at the same time and viewing geometry with respect to Fermi , with comparable detector background rates, we /uniFB01 nd that GRB /uniF0A0 170817A could have been ∼ 30% dimmer before falling below the on-board triggering threshold ( Goldstein et al. 2017 ) , corresponding to a maximum detection distance of about 50 Mpc. An approximate measure of the detectability distance given optimal detection conditions ( e.g., low background, good geometry ) suggests the maximum distance we could have detected this burst is about 80 Mpc -closer than any  \nother SGRB with a /uniFB01 rmly determined redshift. While the GW horizon has been considered the limiting factor for joint detections with EM signals, this joint detection shows that we now must also account for an SGRB detection horizon given the sensitivity of the current gamma-ray observatories.  \nIn addition to limited gamma-ray detector sensitivity, determining the redshift from EM observations alone is more dif /uniFB01 cult for SGRBs than for long GRBs. The fraction of SGRBs with detected X-ray afterglows for Swift BAT detected bursts is ∼ 75% ( Fong et al. 2015 ) , compared to over 90% for long GRBs. 185 It is possible that SGRBs with subluminous prompt gamma-ray emission also have correspondingly weaker X-ray afterglows, and these could account for a large fraction of the quarter of SGRBs without X-ray detections. Even when the X-ray afterglows are detected, they are fainter and thus fade below detectability threshold faster than the afterglows of long GRBs, making direct measurement of the redshift from the afterglow exceedingly rare ( Fong et al. 2015 ) . For SGRBs, the redshift is instead usually determined from the host galaxy. This requires /uniFB01 rst that the afterglow be tied to a particular host galaxy, which can be dif /uniFB01 cult because the SGRB progenitors sometimes lie outside their putative hosts, owing to the natal kicks induced by the supernovae that produced the compact objects in the progenitor system ( Wong et al. 2010 ) . A welllocalized ( ∼ few arcseconds ) SGRB afterglow is associated with a galaxy within a small angular distance on the sky, using probabilistic arguments about chance alignment, and then the redshift of the host galaxy is measured.  \nAppendix B lists all SGRBs with possible redshifts. Most of the list was compiled by combining three relatively complete and recent literature samples ( Fong et al. 2015; Lien et al. 2016; Siellez et al. 2016 ) . Nearly all of these were detected by Swift BAT. It has been suggested that the BAT SGRB distribution is contaminated by the short tail of the long GRB distribution ( Bromberg et al. 2013 ) . Burns et al. ( 2016 ) /uniFB01 nd that the BAT sample is not signi /uniFB01 cantly more contaminated than the GBM sample and the redshift distribution based on Swift BAT SGRBs is therefore a valid proxy for the redshift distribution of GBM SGRBs in the following discussion.  \nBerger ( 2010 ) discuss ' hostless ' SGRBs, which are welllocalized SGRBs that have no obvious associated host galaxy despite deep observational limits. They suggested the hosts could be nearby galaxies at larger angular offset to the afterglow than others farther away, but also put forward the possibility of more distant, undetected hosts. Tunnicliffe et al. ( 2014 ) show that hostless SGRBs have an excess of nearby galaxies within a few arcminutes, relative to long GRBs or random positions, suggesting that at least some of these hostless SGRBs have nearby hosts. Therefore, the traditional assignment of probability of an SGRB to a host galaxy based solely on angular offset from a well-localized afterglow may exclude real associations with larger offsets, which are more likely to be measured for nearby events.  \nTunnicliffe et al. ( 2014 ) also includes the closest potential host for an SGRB prior to GRB /uniF0A0 170817A, with 81 Mpc for GRB 111020A. If real, this association implies an extremely low E iso of 10 46 ~ erg, similar to GRB /uniF0A0 170817A. In light of the secure association of GRB /uniF0A0 170817A with GW170817, a subenergetic E iso may no longer be a reason to doubt subluminous nearby SGRB host associations, and may suggest  \na reconsideration of very nearby host galaxies with large projected angular offsets for hostless SGRBs. We include these putative host associations in Table 2. Also included are SGRBs with extended emission and bursts that have durations exceeding the standard T 90 cut but are believed to be short based on other evidence such as spectral hardness. Asterisks indicate bursts, where the Swift BAT T 2s 90 < , that have localizations of a few arcseconds or better ( as larger localizations increase the chance of false associations due to chance alignment ) , and for which the angular offset of the afterglow from the host ful /uniFB01 lls standard association criteria. For further analysis this restricted sample is our ' gold sample, ' and the full sample the ' total sample. '  \nOne outstanding question is why we have not detected other SGRBs as close as GRB /uniF0A0 170817A. We have established that bursts as dim as GRB /uniF0A0 170817A will not be detected by current gamma-ray instruments if they lie much farther away than GRB /uniF0A0 170817A. This raises questions about GRBs with luminosities between GRB /uniF0A0 170817A and the rest of the GRBs with known luminosities. Some of these are surely being detected, albeit with unassigned redshifts and thus luminosities. While there are only ∼ 40 SGRBs with possible redshifts, several hundreds have been detected without an assigned redshift. Nearby, subluminous SGRBs surely lie among them. There is a lower priority for following-up weak SGRBs, so if nearby events are systematically detected as weak bursts they may not have the required follow-up observations at lower wavelengths to determine the distance to the burst. It could also be that these weak bursts also have lower brightness at lower wavelengths, making them harder to detect even with followup observations. Lastly, we could be detecting these bursts in gamma-rays and X-rays, but failing to properly associate them with their hosts as discussed above.  \nGRB /uniF0A0 170817A is unique in that its distance was /uniFB01 rst measured by GWs, which are currently detectable out to limited distances ( roughly 100 Mpc ) compared to other SGRBs with known redshifts ( see Table 2 ) . This is analogous to the /uniFB01 rst association of long GRBs with supernovae. Long GRBs have redshifts systematically higher than SGRBs ( Coward et al. 2013 ) . The long burst GRB 980425 is the closest GRB to date with a measured distance ( and the only GRB of any class closer than GRB /uniF0A0 170817A ) , and it was the /uniFB01 rst long GRB associated with a supernovae. GRB 980425 was 4 orders of magnitude less energetic than other GRBs detected at that time ( Galama et al. 1998 ) . Because supernovae are less luminous than long GRBs, the long GRBs that are associated with supernovae are systematically closer than the average population. Because of the Malmquist bias, a bias toward detecting intrinsically bright objects ( Malmquist 1922 ) , we only see dim long GRBs when they are nearby. This explains the subluminous nature of GRB 980425, and this observational peculiarity has been con /uniFB01 rmed by other subluminous long GRB-SN detections, including GRB 031203 / SN2003lw ( Malesani et al. 2004 ) , GRB 060218 / SN2006aj ( Modjaz et al. 2006 ) , and GRB 100316D / SN2010bh ( Cano et al. 2011 ) , quanti /uniFB01 ed as a population in Howell & Coward ( 2012 ) . The history of GRB 980425, the other nearby subluminous long GRBs associated with supernova, and the lack of correlation between SGRB gamma-ray /uniFB02 uence and redshift noted by Burns et al. ( 2016 ) motivates the further development of subthreshold searches for counterparts to GW events and for subthreshold SGRBs in general. While GRB /uniF0A0 170817A occurred nearby, and its favorable geometry  \nto Fermi resulted in an on-board trigger, we anticipate that these untriggered searches of GBM and other gamma-ray data will uncover future counterparts to this GW-selected SGRB population.  \nSuggestions of nearby subluminous SGRB populations existed prior to this discovery ( Tanvir et al. 2005; Siellez et al. 2016 ) . Tanvir et al. ( 2005 ) /uniFB01 nd a statistically signi /uniFB01 cant correlation between a large sample of coarsely localized SGRBs detected by the Burst And Transient Source Experiment ( BATSE ) and a sample of nearby galaxies. Without associating individual SGRBs with a potential host, they conclude that ∼ 10% of the SGRB sample could be part of a nearby subluminous population. Siellez et al. ( 2016 ) infer the presence of a nearby subluminous SGRB population through a study of SGRBs with known redshift in the context of BNS and NS -BH population evolution. They /uniFB01 nd an excess of actual nearby low-luminosity SGRBs using the results from their simulations, covering a broad range of assumed lifetimes for the binary system prior to merger.  \nGiant /uniFB02 ares from the highly magnetized NSs known as magnetars can be detected outside our galaxy, with the sole extragalactic example tied to its host coming from SGR 052566, in the Large Magellanic Cloud ( Evans et al. 1980 ) . A giant /uniFB02 are from the galactic magnetar SGR 1806-20 showed a gamma-ray spectrum measured by the Konus -Wind instrument that was well/uniFB01 t by a blackbody with temperature ∼ 175 keV ( Hurley et al. 2005 ) , harder than a regular magnetar burst. This hard spectrum led to the idea that giant /uniFB02 ares from magnetars in nearby galaxies might be a sub-population hiding among the general SGRB population ( Hurley et al. 2005 ) . Tanvir et al. ( 2005 ) found a stronger correlation of BATSE SGRBs with early-type than late-type galaxies, which is not expected if nearby SGRBs arise from giant magnetar /uniFB02 ares in nearby galaxies, but is consistent with a BNS origin. GRB /uniF0A0 170817A is clearly associated with a BNS merger, but even without the connection to GW170817, the spectrum of GRB /uniF0A0 170817A in the GBM data strongly disfavored the BB /uniFB01 t expected for a giant magnetar /uniFB02 are. Another possible signature of a giant magnetar /uniFB02 are is the ringing in its tail at the NS period of a few seconds, which could be detected by GBM or by SPI-ACS for /uniFB02 ares outside our galaxy providing it was close enough. A search for periodic or quasi-periodic emission in the GBM data for GRB /uniF0A0 170817A ( Goldstein et al. 2017 ) found no periodic modulation, providing another discriminant between SGRBs and nearby extragalactic giant magnetar /uniFB02 ares that might be masquerading as SGRBs.", '6.4. Predicted Detection Rates': "The intrinsic speci /uniFB01 c volumetric SGRB rate is often quoted to be around 10 Gpc yr 3 1 --( see, e.g., Guetta & Piran 2006; Coward et al. 2012; Fong et al. 2015; with the value originating from Nakar et al. 2006, who noted that the true rate could be much higher ) . However, unlike GW signals, SGRBs do not have a clear relationship between the observed distance and brightness. As discussed in the previous section, this can be due to intrinsic variations in SGRB luminosities, as well as structure in the jet. In this Section, we investigate the former scenario presenting the implications of GW170817 / GRB /uniF0A0 170817A for future GW and SGRB observations in terms of a simple standard model for the SGRB luminosity distribution. Similar interpretations for other, perhaps more elaborate, models are straightforward.  \nWe model the SGRB luminosity function as a broken power law, with a logarithmic distribution 186  \nL L L L L L L L L , 20 o iso iso iso iso iso L L     f = < > a b --/uni23A7 /uni23A8 /uni23AA /uni23AA /uni23A9 /uni23AA /uni23AA /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F ( ) ( )  \nwhere L iso is the peak isotropic luminosity ( in the source frame ) between 1 keV and 10 MeV, and L a and L b give the power law decay below and above the break at L å . 187 Here, we follow Wanderman & Piran ( 2015 ) in using L 2 10 52  '  erg s -1 , 1 L a  , and 2 L b  . The other important parameter is the minimum SGRB luminosity, which determines the lower cutoff of the luminosity distribution. This is poorly constrained as only nearby low luminosity SGRBs are observable. In Wanderman & Piran ( 2015 ) the minimum luminosity is taken to be L 5 10 min 49 = ' erg s -1 , while other studies use values ranging from 1 10 49 ' erg s -1 to few 10 50 ' erg s -1 ( Regimbau et al. 2015 ) . We assume a threshold value for detectability in GBM of 2 photons cm -2 s -1 for the 64 ms peak photon /uniFB02 ux in the 50 -300 keV band, which is higher than the minimum detectability value to account for the sky-dependent sensitivity of GBM. Furthermore, we model the SGRB spectrum using the Band function with parameters taken from Wanderman & Piran ( 2015 )( namely, E 800 peak = keV, 0.5 Band a =-,and Band b = 2.25 -) . This spectrum is signi /uniFB01 cantly harder than the one observed for GRB /uniF0A0 170817A. The cumulative observed rate predicted for GBM by this base model is shown as a function of redshift in Figure 6 by the purple solid curve.  \nAs discussed in Section 6.1,theinferred L iso is 1.6 0.6  ' ( ) 10 47 erg s -1 , which is signi /uniFB01 cantly lower than any previously detected SGRB, and thus is in tension with this model. In particular, we must extend the lower limit of the luminosity down by a factor of at least 500. At present, there is rather little information available about the low luminosity distribution due to the observational biases discussed in Section 6.3 and, consequently, there is a signi /uniFB01 cant degeneracy between the minimum SGRB luminosity and the rate ( Wanderman & Piran 2015 ) .Letus consider the most straightforward extension of the above model and set L 1 10 min 47 = ' erg s -1 while maintaining 1 L a = .In order to retain the same prediction for high-luminosity SGRBs, this requires a 500-fold increase in the number of SGRBs, with the majority emitting at low luminosity. The cumulative observed rate predicted for GBM by this simple extension is shown as a function of redshift in Figure 6 by the red solid curve, and is comparable to the measured BNS merger rate shown in black. This simple extension would imply SGRBs are not beamed and that essentially all BNS mergers are accompanied by at least a subluminous SGRB.  \nTherefore, to reduce this tension and explore other possible extensions, we introduce an additional power law break below  \nFigure 6. Predicted detection rates per year as a function of redshift. The red, blue, and green solid lines refer to the GBM observed SGRB rate assuming a minimum luminosity L min of 1 10 47 ' erg s -1 , and 1 L a = , 2 L b = and 1, 0.5, 0 L g = { } in Equation ( 21 ) , respectively. The purple solid line refers to the base model with L min of 5 10 49 ' erg s -1 . The four curves are normalized by imposing 40 triggered SGRB per year. As L g increases, the observed rate is no longer volumetric at lower and lower redshifts, because a fraction of SGRBs becomes too dim to be detected. For reference, the red, blue and green dot-dashed curves show the local SGRB occurrence rate for L 1 10 min 47 = ' erg s -1 and 1, 0.5, 0 L g = { } , respectively. The black line and gray band show the BNS merger rate 1540 Gpc yr 1220 3200 3 1 -+ --determined with the detection of GW170817 ( Abbott et al. 2017e ) . For comparison, the measured SGRBs redshift distribution from Table 2 is shown in cyan, and is broadly compatible with all of the models. The dotted vertical cyan line refers to the redshift of GRB /uniF0A0 170817A host galaxy.  \n<!-- image -->  \nL 5 10 49  = ' erg s -1 :  \nL L L L L L L L L L L L L L L L , 21 o iso iso iso iso iso iso iso L L L L          f = < < < > g a a b ----/uni23A7 /uni23A8 /uni23AA /uni23AA /uni23AA /uni23AA /uni23A9 /uni23AA /uni23AA /uni23AA /uni23AA /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F ( ) ( )  \nand consider three values 1, 0.5, 0 L g = { } for the power law index below this second break. We normalize these three cases to 40 triggered SGRBs per year for GBM, and the 1 L g = case corresponds to the simple extension discussed above. For reference, Figure 6 shows the local SGRB occurrence rate for L 1 10 min 47 = ' erg s -1 and 1, 0.5, 0 L g = { } indicated by the red, blue, and green dotted curves, respectively, and the BNS merger rate 1540 Gpc yr 1220 3200 3 1 -+ --determined with the detection of GW170817 ( gray band, with the mean in black; Abbott et al. 2017e ) . The 1 L g = case produces the largest number of subluminous SGRBs, and leads to a sharp departure at redshift ∼ 0.005 from volumetric detection to detection limited by the GBM sensitivity. For 0 L g = the transition is smoother as there is only a small number of subluminous SGRBs, and the observed rate departs gradually from the occurrence rate.  \nWhen we include the luminosity of GRB /uniF0A0 170817A by setting L 1 10 min 47 = ' erg s -1 and 1 L g = , the expected  \ndetection rate at a redshift of z 0.1 » is around a factor of 2 higher than for the Wanderman & Piran ( 2015 ) model. At a redshift of z 0.01 » , which is close to the observed redshift for GRB /uniF0A0 170817A, rather than expecting to observe 1 event per 650 years with GBM, this is increased to 1 per year. The expected detection rate at a this redshift for the 0.5 L g = and 0 L g = extensions is of roughly 1 observed event per 10 and 65 years, respectively. The expectations we obtain for GBM are consistent with the distribution of SGRBs with known redshifts reported in Table 2, in Appendix B.  \nUsing the BNS merger volumetric rate estimated from GW170817 as a new input to the detection rate calculation presented in Abbott et al. ( 2017a ) , the LIGO -Virgo detection rate is narrowed down from 0.04 -100 to ∼ 1 -50 BNS coalescences during the 2018 -19 observing run, with the remaining uncertainty arising in part from the not-yet-known detector sensitivities during that run. At design sensitivity, the LIGO and Virgo detectors can expect to detect ∼ 6 -120 BNS coalescences per year, as opposed to the previously estimated 0.1 -200 BNS coalescences per year. Inclusion of any additional BNS detections in the meantime will allow this prediction to be further sharpened.  \nIndependently, we use the GBM detection rate as a function of redshift to predict joint GW -GRB BNS detection rates ( Clark et al. 2015 ) . Both the rates and their relative uncertainties are signi /uniFB01 cantly reduced, compared to the GWonly detection rate estimates above, since the majority of distant mergers will be undetectable by GBM and the GBM  \nSGRB detection rate is well measured. The degree to which GBM-selected SGRBs are preferentially on-axis is unclear. When estimating the joint detection rate we include both no selection and a pure on-axis selection hypothesis, the latter implying a larger detection probability by the LIGO and Virgo detectors. During the 2018 -19 observing run, we expect 0.1 -1.4 joint detections per year for the GW interferometer network and GBM triggered SGRBs, with the high end of the interval corresponding to the 1 L a = , L 1 10 min 47 = ' erg s -1 extension of the luminosity function. At design sensitivity, the expected joint detection rate increases to 0.3 -1.7 per year.  \nFuture joint GW-SGRB observations will provide signi /uniFB01 cant new insights into low-luminosity SGRBs. In particular, both joint observations and lower limits on distances to SGRBs not observed in GWs ( e.g., GRB 150906B in Abbott et al. 2017b ) will constrain the rate of nearby GRBs. Future GW observations of BNS mergers will reduce the uncertainty in the rate of such events, while observation of GW signals with no SGRB counterpart will limit the SGRB beaming angles. Finally, subthreshold searches in GRB data around the time of GW events could signi /uniFB01 cantly increase the number of joint observations.", '6.5. Limits on Precursor and Extended Emission': "At gamma-ray energies, SGRBs are characterized by a prompt emission episode lasting at most ∼ 2 s. Observational evidence for precursor /uniFB02 ares associated with SGRBs ( Troja et al. 2010; Burns 2017; Minaev & Pozanenko 2017 ) and temporally extended emission ( Lazzati et al. 2001; Connaughton 2002 ) is so far inconclusive. Given the small distance to the source, the absence of such emission from GRB /uniF0A0 170817A provides an important data point and may constrain models that predict it. The /uniFB02 ux upper limits set in Section 2.2 correspond to an intrinsic upper limit of 2.4 10 47 ~ ' erg s -1 for precursor emission on the 0.1 s timescale, 7.0 10 46 ~ ' erg s -1 for precursor emission on the 1.0 s timescale, and 2.2 10 46 ~ ' erg s -1 for extended emission on the 10 s timescale.  \nMagnetospheric interactions in NS binaries have been proposed as a source of nearly isotropic emission preceding the merger ( e.g., Hansen & Lyutikov 2001; Metzger & Zivancev 2016 ) . In the context of such models, the nondetection of precursors associated with GW170817 suggests the absence of strong magnetic /uniFB01 elds in the last ∼ 200 s before merger. Hansen & Lyutikov ( 2001 ) , for instance, predict a luminosity that depends on the magnetic /uniFB01 eld B as L B a 7.4 10 10 G 10 cm erg s 45 15 2 7 7 1 ' -- ( ) ( ) . We can combine this estimate with the least-constraining GBM intrinsic upper limit above and assume a /uniFB01 nal separation a 3 10 cm 6 = ' before disruption. The resulting limit is B 8 10 13 < ' G, which is weaker than the magnetic /uniFB01 elds of most known magnetars ( Olausen & Kaspi 2014 ) .However, the GBM upper limit still lies within the luminosity range of other similar models ( Metzger & Zivancev 2016 ) .  \nResonant shattering of the NS crust has also been proposed as a source of emission prior to merger, with a maximum time delay of tens of seconds and nearly isotropic angular distribution ( Tsang et al. 2012 ) . The luminosity of such precursor emission depends on the crust breaking strain b  and the emission timescale t D as L t 7 10 48 b 2  ' D  erg, from which we can derive t 10 1 s b 2 2   D -( ) . Assuming this mechanism took place in GW170817, and taking 0.1 b  =  \n( Horowitz & Kadau 2009 ) , the emission either lasted more than a few seconds or happened below the GBM energy range, i.e., ∼ 10 keV. There might also be a dependence of the luminosity on the details of the NS EOS, although that is yet to be investigated in detail. Similarly to magnetospheric interaction, however, resonant shattering emission ultimately requires a suf /uniFB01 ciently large magnetic /uniFB01 eld and a simple explanation for the absence of a signal is again the lack of intense magnetic /uniFB01 elds prior to merger.  \nGBM and SPI-ACS observed no temporally extended gamma-ray emission for GRB /uniF0A0 170817A. Such emission would be a signature of a long-lived NS remnant powering the SGRB and our /uniFB02 ux limits may suggest instead that the remnant is a BH. Metzger et al. ( 2008 ) invoke a long-lived millisecond magnetar to explain SGRBs with extended gamma-ray emission ( Norris & Bonnell 2006 ) , and millisecond magnetars have also been suggested as possible causes for the plateaus seen in X-ray afterglows of some SGRBs ( Rowlinson et al. 2013 ) . The earliest X-ray observation was only performed 50 ks after GRB /uniF0A0 170817A ( Evans et al. 2017 ) and hence limits are only set after this time. Future observations may further constrain this scenario, e.g., radio observations on the timescale of a year ( Fong et al. 2016 ) .  \nWe encourage the development of quantitative predictions of luminosity as a function of energy, time and physical parameters of the source, as the multiple upcoming joint observations of BNS mergers suggest the possibility of interesting constraints on the pre-merger physics.", '7. Conclusion': "The joint observation of GW170817 and GRB /uniF0A0 170817A con /uniFB01 rms the association of SGRBs with BNS mergers. With just one joint event, we have set stringent limits on fundamental physics and probed the central engine of SGRBs in ways that have not been possible with EM data alone, demonstrating the importance of multi-messenger astronomy.  \nThe small time offset and independent localizations, though coarse, allowed an unambiguous association of these two events. Because GRB /uniF0A0 170817A occurred nearby, an autonomous trigger on-board GBM alerted follow-up observers to the presence of a counterpart to GW170817. At design sensitivity, however, Advanced LIGO and Virgo could in principle detect GW170817 beyond the distance that any active gamma-ray observatory would trigger on a burst like GRB /uniF0A0 170817A. Subthreshold searches for SGRBs can extend the gamma-ray horizon and the detection of GRB /uniF0A0 170817A provides motivation for further subthreshold search development.  \nA joint detection at greater distance and for an SGRB with more typical energetics would allow tighter constraints on the temporal offset and the derived inferences. Should NS -BH binaries also be SGRB progenitors, only a joint detection between GW and EM can provide decisive evidence.  \nIn this Letter we propose several explanations for the observed dimness of GRB /uniF0A0 170817A. We suggest joint detections should be more common than previously predicted, and future observations of multiple events should enable a study of the populations of mergers and their associated SGRBs, shedding light on the jet geometry and intrinsic brightness distribution. Furthermore, detections with multiple  \nGW interferometers can provide more stringent constraints on the inclination angles of these systems. The joint detections of SGRBs arising from BNS and NS -BH mergers will constrain the fraction of SGRBs originating from each progenitor class.  \nThe global network of GW detectors and wide/uniFB01 eld gammaray instruments, such as Fermi -GBM and INTEGRAL / SPIACS, are critical to the future of multi-messenger astronomy in the GW era.  \nWe dedicate this Letter to the memory of Neil Gehrels. His pioneering work in gamma-ray astronomy and his vision for multi-messenger astrophysics were instrumental to our discoveries.  \nThe authors thank the referees for their invaluable comments and feedback, especially in a timely manner. The authors gratefully acknowledge the support of the United States National Science Foundation ( NSF ) for the construction and operation of the LIGO Laboratory and Advanced LIGO as well as the Science and Technology Facilities Council ( STFC ) of the United Kingdom, the Max-PlanckSociety ( MPS ) , and the State of Niedersachsen / Germany for support of the construction of Advanced LIGO and construction and operation of the GEO600 detector. Additional support for Advanced LIGO was provided by the Australian Research Council. The authors gratefully acknowledge the Italian Istituto Nazionale di Fisica Nucleare ( INFN ) , the French Centre National de la Recherche Scienti /uniFB01 que ( CNRS ) and the Foundation for Fundamental Research on Matter supported by the Netherlands Organisation for Scienti /uniFB01 c Research, for the construction and operation of the Virgo detector and the creation and support of the EGO consortium. The authors also gratefully acknowledge research support from these agencies as well as by the Council of Scienti /uniFB01 cand Industrial Research of India, the Department of Science and Technology, India, the Science & Engineering Research Board ( SERB ) , India, the Ministry of Human Resource Development, India, the Spanish Agencia Estatal de Investigación, the Vicepresidència i Conselleria d ' Innovació Recerca i Turisme and the Conselleria d ' Educació i Universitat del Govern de les Illes Balears, the Conselleria d ' Educació Investigació Cultura i Esport de la Generalitat Valenciana, the National Science Centre of Poland, the Swiss National Science Foundation ( SNSF ) , the Russian Foundation for Basic Research, the Russian Science Foundation, the European Commission, the European Regional Development Funds ( ERDF ) , the Royal Society, the Scottish Funding Council, the Scottish Universities Physics Alliance, the Hungarian Scienti /uniFB01 c Research Fund ( OTKA ) , the Lyon Institute of Origins ( LIO ) , the National Research, DevelopmentandInnovationOf /uniFB01 ce Hungary ( NKFI ) , the National Research Foundation of Korea, Industry Canada and the Province of Ontario through the Ministry of Economic Development and Innovation, the Natural Science and Engineering Research Council Canada, the Canadian Institute for Advanced Research, the Brazilian Ministry of Science, Technology, Innovations, and Communications, the International Center for Theoretical Physics South American Institute for Fundamental Research ( ICTP-SAIFR ) , the Research Grants Council of Hong Kong, the National Natural Science Foundation of China ( NSFC ) , the Leverhulme Trust, the  \nResearch Corporation, the Ministry of Science and Technology ( MOST ) , Taiwan and the Kavli Foundation. The authors gratefully acknowledge the support of the NSF, STFC, MPS, INFN, CNRS and the State of Niedersachsen / Germany for provision of computational resources.  \nThe USRA co-authors gratefully acknowledge NASA funding through contract NNM13AA43C. The UAH coauthors gratefully acknowledge NASA funding from cooperative agreement NNM11AA01A. E.B. and T.D.C. are supported by an appointment to the NASA Postdoctoral Program at the Goddard Space Flight Center, administered by Universities Space Research Association under contract with NASA. D.K., C.A.W.H., C.M.H., and T.L. gratefully acknowledge NASA funding through the Fermi GBM project. Support for the German contribution to GBM was provided by the Bundesministerium für Bildung und Forschung ( BMBF ) via the Deutsches Zentrum für Luft und Raumfahrt ( DLR ) under contract number 50 QV 0301. A.v.K. was supported by the Bundesministeriums fr Wirtschaft und Technologie ( BMWi ) through DLR grant 50 OG 1101. N.C. and J.B. acknowledge support from NSF under grant PHY-1505373. S.M.B. acknowledges support from Science Foundation Ireland under grant 12 / IP / 1288.  \nThis work is based on observations with INTEGRAL ,an ESA project with instruments and science data center funded by ESA member states ( especially the PI countries: Denmark, France, Germany, Italy, Switzerland, Spain ) , and with the participation of Russia and the USA. The INTEGRAL SPI project has been completed under the responsibility and leadership of CNES. The SPI-ACS detector system has been provided by MPE Garching / Germany. The SPI team is grateful to ASI, CEA, CNES, DLR, ESA, INTA, NASA and OSTC for their support. The Italian INTEGRAL team acknowledges the support of ASI / INAF agreement n. 2013025-R.1. R.D. and A.v.K. acknowledge the German INTEGRAL support through DLR grant 50 OG 1101. A.L. and R.S. acknowledge the support from the Russian Science Foundation ( grant 14-22-00271 ) . A.D. is funded by Spanish MINECO / FEDER grant ESP2015-65712-C5-1-R. Some of the results in this paper have been derived using the HEALPix ( Gorski et al. 2005 ) package. We are grateful VirtualData from LABEX P2IO for enabling access to the StratusLab academic cloud. We acknowledge the continuous support by the INTEGRAL Users Group and the exceptionally ef /uniFB01 cient support by the teams at ESAC and ESOC for the scheduling of the targeted follow-up observations.", 'Appendix A Full Derivation of Photospheric Radius': "In order to explain the observed soft BB component, we present a model that simultaneously yields predictions on ( 1 ) the time difference between the GW emission and the beginning of the gamma-ray radiation, ( 2 ) the estimated temperature of the BB component ( k T 10.3 1.5 B BB obs =  ( ) keV ) and ( 3 ) its average isotropic luminosity, L E t 1.1 iso,BB iso,BB BB = D = ' 10 46 erg s -1 , where we have used the source interval width of the soft tail spectral /uniFB01 tof t 1.15 BB D = s.  \nFollowing MR00, our model depends upon three main parameters to provide predictions of the aforementioned three observable quantities. First, the radius from which the /uniFB01 reball  \nis initiated, which is assumed to be R GM c 6 0 BH 2 = ( the innermost stable circular orbit of a Schwarzschild BH with a mass equal to M BH ) .Inourcase R 2.5 10 0 6 = ' cm, and M M 2.8 BH =  . From numerical models of GRB jets produced in BNS merger remnants ( e.g., Aloy et al. 2005 ) , the value of R 0 can be associated to the stagnation point of a relativistic out /uniFB02 ow, and it is fairly well constrained to be a few gravitational radii of the BH. Second, the initial luminosity of the /uniFB01 reball, L 0. This is a free parameter of themodelandwenoteitcanbemuchlargerthan L iso,BB , since the /uniFB01 reball must expand from its initial volume ( R 0 3 ~ ) to the size where the photosphere appears ( see below ) .Thisis also needed, since the observed gamma-ray luminosity, will be a fraction 0 1 r   < of the total ( kinetic ) luminosity. Finally, the third parameter is the dimensionless entropy of the /uniFB01 reball, η .Fora /uniFB01 reball baryon load M ˙ and a luminosity L 0, L Mc 2 h = ( ˙ ) . Typical values of η are larger than 100 to prevent the compactness problem ( Goodman 1986 ) ,whichis not an issue for GRB /uniF0A0 170817A owing to the lack of emission detected above 511 keV. Therefore, our model may allow for values of η substantially smaller than ∼ 100. As we shall see, a combination of L 10 0 50  erg and 18 h  , results in a viable model to account for the delay of GRB with respect to the GW signal and the average luminosity of the soft BB component.  \nThe initial BB temperature in units of the electron rest mass is given by ( MR00, Equation ( 5 ))  \nk mc L R ca L R 4 1.5 10 erg s 2.5 10 cm , 22 e r 0 B 2 0 0 2 1 4 0 50 1 1 4 0 6 1 2 p Q= ' -- /uni239C /uni239F /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F /uni239B /uni239D /uni239E /uni23A0 ( )  \nwhere a 7.57 10 r 18 = ' -kg cm -1 s -2 K -4 is the radiation constant. The value of 0 Q corresponds to a comoving temperature k T 750 B 0  keV. The radius at which the internal energy of the /uniFB01 reball is converted into kinetic energy, i.e., the saturation radius is  \nR R R 4.5 10 cm 18 2.5 10 cm . 23 s 0 7 0 6 h h = ' '  /uni239C /uni239F /uni239C /uni239F /uni239B /uni239D /uni239E /uni23A0 /uni239B /uni239D /uni239E /uni23A0 ( )  \nThe critical baryon load * h below which the photosphere of an expanding /uniFB01 reball happens after the /uniFB01 reball coasts at constant Lorentz factor g h  , i.e., at radii larger than R s , is given by ( MR00 )  \nL Y mcR 4 , 24 T 0 p 3 0 1 4 * h s p = /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F ( )  \nwhere Y is the number of electrons per baryon, 6.65 10 T 25 s = ' -cm 2 is the Thomson cross-section. Using /uniFB01 ducial values, we obtain  \nL R Y 470 10 erg s 2.5 10 cm . 25 0 50 1 1 4 0 6 1 4 1 4 * h ' -- /uni239C /uni239F /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F /uni239B /uni239D /uni239E /uni23A0 ( )  \nIn the previous expression, we have taken Y 1  , which is appropriate once pairs are not present in the system. This is the case for radii larger than R p ( MR00 )  \nR R L R 1.1 10 cm 10 erg s 2.5 10 cm 0.03 , 26 p p p 0 0 8 0 50 1 1 4 0 6 1 2 1 = Q Q ' ' ' Q -- /uni239C /uni239F /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F /uni239B /uni239D /uni239E /uni23A0 /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F ( )  \nwhere the comoving dimensionless temperature below which e  pairs drop out of equilibrium is 0.03 17 p Q   keV. Since we have set * h h < , the photosphere will happen at a radius ( MR00 )  \nR L Y mc L Y 4 2 10 cm 10 erg s 18 . 27 T ph p 3 3 13 0 50 1 3 s p h h = ' -- /uni239C /uni239F /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F /uni239B /uni239D /uni239E /uni23A0 ( )  \nNote that R R p ph  for our choice of tunable parameters.  \nThe laboratory frame time needed for the /uniFB01 reball expanding at roughly the speed of light to reach the transparency radius is t R c ph ph  , thus,  \nt L Y 672 s 10 erg s 18 . 28 ph 0 50 1 3 h -- /uni239C /uni239F /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F /uni239B /uni239D /uni239E /uni23A0 ( )  \nTo compute the time delay between a photon emitted at R 0 at t /uniF0A0 = /uniF0A0 0 ( namely, signaling the GW detection ) and another one at R ph , we must consider that the /uniFB01 reball begins its expansion from rest, in which case we shall apply the following relation between the arrival time and the time at which the photosphere appears  \nt t 1 2 1 cos , 29 a ph 2 h J + - /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F ( )  \nwhere ϑ is the angle between the radial direction and the line of sight and cos J takes values only in the interval vc ,1 [ ] ( Bianco et al. 2001 ) . Using /uniFB01 ducial values for η , we obtain  \nt t L Y 2 1.0 s 10 erg s 18 , for cos 1 a 1 ph 2 0 50 1 5 g h J = --  /uni239C /uni239F /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F /uni239B /uni239D /uni239E /uni23A0 ( )  \n(  \n)  \n30  \n31 t t L Y v c 2.1 s 10 erg s 18 , for cos . a 2 ph 2 0 50 1 5 g h J = --  /uni239C /uni239F /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F /uni239B /uni239D /uni239E /uni23A0 ( ) ( )  \nThese values account for a signi /uniFB01 cant fraction of the time delay between the GW and the prompt radiation in the soft thermal peak.  \nWe can now estimate the comoving temperature of the expanding /uniFB01 reball at the photospheric radius ( MR00 )  \nkT kT R R 32 B ph B 0 ph s 2 3 = -/uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F ( )  \nL R Y 0.13 keV 10 erg s 2.5 10 cm 18 . 33 0 50 1 5 12 0 6 5 6 8 3 2 3 h ' ' ---- /uni239C /uni239F /uni239C /uni239F /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F /uni239B /uni239D /uni239E /uni23A0 /uni239B /uni239D /uni239E /uni23A0 ( )  \nThis corresponds to an observed temperature, T T ph obs ph h =  \nkT kT L R Y 2.3 keV 10 erg s 2.5 10 cm 18 . 34 B ph obs B ph 0 50 1 5 12 0 6 5 6 11 3 2 3 h h = ' ' ---- /uni239C /uni239F /uni239C /uni239F /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F /uni239B /uni239D /uni239E /uni23A0 /uni239B /uni239D /uni239E /uni23A0 ( )  \nThe corresponding BB luminosity is ( MR00 ) ,  \nL L R R 35 ph 0 ph s 2 3 = -/uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F ( )  \nL R Y 1.7 10 erg s 10 erg s 2.5 10 cm 18 . 36 46 1 0 50 1 7 12 0 6 5 6 8 3 2 3 h ' ' ' ---- /uni239C /uni239F /uni239C /uni239F /uni239B /uni239D /uni239C /uni239E /uni23A0 /uni239F /uni239B /uni239D /uni239E /uni23A0 /uni239B /uni239D /uni239E /uni23A0 ( )  \nThus, with the reference values of our model for L 0 and η ,we obtain L L ph iso,BB  . On the other hand, T ph obs underpredicts T BB obs by a factor ∼ 4. Nonetheless, we are neglecting Comptonization effects, which may slightly raise the estimated temperature.", 'Appendix B List of SGRBs with Associated Redshift': "Table 2 is a list of possible redshifts for GRBs that have been argued to belong to the short class. The asterisks show the ' gold sample ' selection with standard cuts on duration ( T 2 90 < s ) and localization uncertainty ( ∼ arcsecond or better ) . The others include SGRBs with extended gamma-ray emission, those slightly longer than 2 s that are spectrally hard or show negligible spectral lag, and bursts that are best localized by Swift BAT ( so a chance association is possible ) . Most of these redshifts come from Lien et al. ( 2016 ) ; Siellez et al. ( 2016 ) , and Fong et al. ( 2015 ) ;forthe original citations see references therein. They also include bursts from Tunnicliffe et al. ( 2014 ) . 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Bisht 22 , M. Bitossi 24,31 , C. Biwer 45 , M. A. Bizouard 29 , J. K. Blackburn 1 , J. Blackman 49 , C. D. Blair 1,65 , D. G. Blair 65 , R. M. Blair 48 , S. Bloemen 66 , O. Bock 10 , N. Bode 10 , M. Boer 67 , G. Bogaert 67 , A. Bohe 39 , F. Bondu 68 , E. Bonilla 52 , R. Bonnand 8 , B. A. Boom 14 , R. Bork 1 , V. Boschi 24,31 , S. Bose 19,69 , K. Bossie 7 , Y. Bouffanais 40 , A. Bozzi 31 , C. Bradaschia 24 , P. R. Brady 21 , M. Branchesi 17,18 , J. E. Brau 70 , T. Briant 71 , A. Brillet 67 , M. Brinkmann 10 , V. Brisson 29 , P. Brockill 21 , J. E. Broida 72 , A. F. Brooks 1 , D. A. Brown 45 , D. D. Brown 73 , S. Brunett 1 , C. C. Buchanan 2 , A. Buikema 15 , T. Bulik 74 , H. J. Bulten 14,75 , A. Buonanno 39,76 , D. Buskulic 8 , C. Buy 40 , R. L. Byer 52 , M. Cabero 10 , L. Cadonati 77 , G. Cagnoli 27,78 , C. Cahillane 1 , J. Calderón Bustillo 77 , T. A. Callister 1 , E. Calloni 4,79 , J. B. Camp 80 , M. Canepa 60,81 , P. Canizares 66 , K. C. Cannon 82 , H. Cao 73 , J. Cao 83 , C. D. Capano 10 , E. Capocasa 40 , F. Carbognani 31 , S. Caride 84 , M. F. Carney 85 , J. Casanueva Diaz 29 , C. Casentini 33,34 , S. Caudill 14,21 , M. Cavaglià 11 , F. Cavalier 29 , R. Cavalieri 31 , G. Cella 24 , C. B. Cepeda 1 , P. Cerdá-Durán 25 , G. Cerretani 23,24 , E. Cesarini 34,86 , S. J. Chamberlin 64 , M. Chan 47 , S. Chao 87 , P. Charlton 88 , E. Chase 89 , E. Chassande-Mottin 40 , D. Chatterjee 21 , K. Chatziioannou 90 , B. D. Cheeseboro 42 , H. Y. Chen 91 , X. Chen 65 , Y. Chen 49 , H.-P. Cheng 5 , H. Chia 5 , A. Chincarini 60 ,  \nA. Chiummo 31 , T. Chmiel 85 , H. S. Cho 92 , M. Cho 76 , J. H. Chow 26 , N. Christensen 67,72 , Q. Chu 65 , A. J. K. Chua 13 , S. Chua 71 , A. K. W. Chung 93 , S. Chung 65 , G. Ciani 5,54,55 , R. Ciol /uniFB01 94,95 , C. E. Cirelli 52 , A. Cirone 60,81 , F. Clara 48 , J. A. Clark 77 , P. Clearwater 96 , F. Cleva 67 , C. Cocchieri 11 , E. Coccia 17,18 , P.-F. Cohadon 71 , D. Cohen 29 , A. Colla 36,97 , C. G. Collette 98 , L. R. Cominsky 99 , M. Constancio Jr. 16 , L. Conti 55 , S. J. Cooper 59 , P. Corban 7 , T. R. Corbitt 2 , I. Cordero-Carrión 100 , K. R. Corley 51 , N. Cornish 101 , A. Corsi 84 , S. Cortese 31 , C. A. Costa 16 , M. W. Coughlin 1,72 , S. B. Coughlin 89 , J.-P. Coulon 67 , S. T. Countryman 51 , P. Couvares 1 , P. B. Covas 102 , E. E. Cowan 77 , D. M. Coward 65 , M. J. Cowart 7 , D. C. Coyne 1 , R. Coyne 84 , J. D. E. Creighton 21 , T. D. Creighton 103 , J. Cripe 2 , S. G. Crowder 104 , T. J. Cullen 2,30 , A. Cumming 47 , L. Cunningham 47 , E. Cuoco 31 , T. Dal Canton 80 , G. Dálya 56 , S. L. Danilishin 10,22 , S. D ' Antonio 34 , K. Danzmann 10,22 , A. Dasgupta 105 , C. F. Da Silva Costa 5 , V. Dattilo 31 , I. Dave 61 , M. Davier 29 , D. Davis 45 , E. J. Daw 106 , B. Day 77 , S. De 45 , D. DeBra 52 , J. Degallaix 27 , M. De Laurentis 4,17 , S. Deléglise 71 , W. Del Pozzo 23,24,59 , N. Demos 15 , T. Denker 10 , T. Dent 10 , R. De Pietri 107,108 , V. Dergachev 39 , R. De Rosa 4,79 , R. T. DeRosa 7 , C. De Rossi 27,31 , R. DeSalvo 109 , O. de Varona 10 , J. Devenson 28 , S. Dhurandhar 19 , M. C. Díaz 103 , L. Di Fiore 4 , M. Di Giovanni 95,110 , T. Di Girolamo 4,51,79 , A. Di Lieto 23,24 , S. Di Pace 36,97 , I. Di Palma 36,97 , F. Di Renzo 23,24 , Z. Doctor 91 , V. Dolique 27 , F. Donovan 15 , K. L. Dooley 11 , S. Doravari 10 , I. Dorrington 37 , R. Douglas 47 , M. Dovale Álvarez 59 , T. P. Downes 21 , M. Drago 10 , C. Dreissigacker 10 , J. C. Driggers 48 , Z. Du 83 , M. Ducrot 8 , P. Dupej 47 , S. E. Dwyer 48 , T. B. Edo 106 , M. C. Edwards 72 , A. Ef /uniFB02 er 7 , H.-B. Eggenstein 10,39 , P. Ehrens 1 , J. Eichholz 1 , S. S. Eikenberry 5 , R. A. Eisenstein 15 , R. C. Essick 15 , D. Estevez 8 , Z. B. Etienne 42 , T. Etzel 1 , M. Evans 15 , T. M. Evans 7 , M. Factourovich 51 , V. Fafone 17,33,34 , H. Fair 45 , S. Fairhurst 37 , X. Fan 83 , S. Farinon 60 , B. Farr 91 , W. M. Farr 59 , E. J. Fauchon-Jones 37 , M. Favata 111 , M. Fays 37 , C. Fee 85 , H. Fehrmann 10 , J. Feicht 1 , M. M. Fejer 52 , A. Fernandez-Galiana 15 , I. Ferrante 23,24 , E. C. Ferreira 16 , F. Ferrini 31 , F. Fidecaro 23,24 , D. Finstad 45 , I. Fiori 31 , D. Fiorucci 40 , M. Fishbach 91 , R. P. Fisher 45 , M. Fitz-Axen 46 , R. Flaminio 27,112 , M. Fletcher 47 , H. Fong 90 , J. A. Font 25,113 , P. W. F. Forsyth 26 , S. S. Forsyth 77 , J.-D. Fournier 67 , S. Frasca 36,97 , F. Frasconi 24 , Z. Frei 56 , A. Freise 59 , R. Frey 70 , V. Frey 29 , E. M. Fries 1 , P. Fritschel 15 , V. V. Frolov 7 , P. Fulda 5 , M. Fyffe 7 , H. Gabbard 47 , B. U. Gadre 19 , S. M. Gaebel 59 , J. R. Gair 114 , L. Gammaitoni 43 , M. R. Ganija 73 , S. G. Gaonkar 19 , C. Garcia-Quiros 102 , F. Garu /uniFB01 4,79 , B. Gateley 48 , S. Gaudio 38 , G. Gaur 115 , V. Gayathri 116 , N. Gehrels 80, † , G. Gemme 60 , E. Genin 31 , A. Gennai 24 , D. George 12 , J. George 61 , L. Gergely 117 , V. Germain 8 , S. Ghonge 77 , Abhirup Ghosh 20 , Archisman Ghosh 14,20 , S. Ghosh 14,21,66 , J. A. Giaime 2,7 , K. D. Giardina 7 , A. Giazotto 24 , K. Gill 38 , L. Glover 109 , E. Goetz 118 , R. Goetz 5 , S. Gomes 37 , B. Goncharov 6 , G. González 2 , J. M. Gonzalez Castro 23,24 , A. Gopakumar 119 , M. L. Gorodetsky 62 , S. E. Gossan 1 , M. Gosselin 31 , R. Gouaty 8 , A. Grado 4,120 , C. Graef 47 , M. Granata 27 , A. Grant 47 , S. Gras 15 , C. Gray 48 , G. Greco 121,122 , A. C. Green 59 , E. M. Gretarsson 38 , P. Groot 66 , H. Grote 10 , S. Grunewald 39 , P. Gruning 29 , G. M. Guidi 121,122 , X. Guo 83 , A. Gupta 64 , M. K. Gupta 105 , K. E. Gushwa 1 , E. K. Gustafson 1 , R. Gustafson 118 , O. Halim 17,18 , B. R. Hall 69 , E. D. Hall 15 , E. Z. Hamilton 37 , G. Hammond 47 , M. Haney 123 , M. M. Hanke 10 , J. Hanks 48 , C. Hanna 64 , M. D. Hannam 37 , O. A. Hannuksela 93 , J. Hanson 7 , T. Hardwick 2 , J. Harms 17,18 , G. M. Harry 124 , I. W. Harry 39 , M. J. Hart 47 , C.-J. Haster 90 , K. Haughian 47 , J. Healy 58 , A. Heidmann 71 , M. C. Heintze 7 , H. Heitmann 67 , P. Hello 29 , G. Hemming 31 , M. Hendry 47 , I. S. Heng 47 , J. Hennig 47 , A. W. Heptonstall 1 , M. Heurs 10,22 , S. Hild 47 , T. Hinderer 66 , D. Hoak 31 , D. Hofman 27 , K. Holt 7 , D. E. Holz 91 , P. Hopkins 37 , C. Horst 21 , J. Hough 47 , E. A. Houston 47 , E. J. Howell 65 , A. Hreibi 67 , Y. M. Hu 10 , E. A. Huerta 12 , D. Huet 29 , B. Hughey 38 , S. Husa 102 , S. H. Huttner 47 , T. Huynh-Dinh 7 , N. Indik 10 , R. Inta 84 , G. Intini 36,97 , H. N. Isa 47 , J.-M. Isac 71 , M. Isi 1 , B. R. Iyer 20 , K. Izumi 48 , T. Jacqmin 71 , K. Jani 77 , P. Jaranowski 125 , S. Jawahar 63 , F. Jiménez-Forteza 102 , W. W. Johnson 2 , N. K. Johnson-McDaniel 13 , D. I. Jones 126 , R. Jones 47 , R. J. G. Jonker 14 , L. Ju 65 , J. Junker 10 , C. V. Kalaghatgi 37 , V. Kalogera 89 , B. Kamai 1 , S. Kandhasamy 7 , G. Kang 41 , J. B. Kanner 1 , S. J. Kapadia 21 , S. Karki 70 , K. S. Karvinen 10 , M. Kasprzack 2 , W. Kastaun 10 , M. Katolik 12 , E. Katsavounidis 15 , W. Katzman 7 , S. Kaufer 22 , K. Kawabe 48 , F. Kéfélian 67 , D. Keitel 47 , A. J. 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Lee 136 , H. W. Lee 130 , K. Lee 47 , J. Lehmann 10 , A. Lenon 42 , M. Leonardi 95,110 , N. Leroy 29 , N. Letendre 8 , Y. Levin 6 , T. G. F. Li 93 , S. D. Linker 109 , T. B. Littenberg 137 , J. Liu 65 , R. K. L. Lo 93 , N. A. Lockerbie 63 , L. T. London 37 , J. E. Lord 45 , M. Lorenzini 17,18 , V. Loriette 138 , M. Lormand 7 , G. Losurdo 24 , J. D. Lough 10 , C. O. Lousto 58 , G. Lovelace 30 , H. Lück 10,22 , D. Lumaca 33,34 , A. P. Lundgren 10 , R. Lynch 15 , Y. Ma 49 , R. Macas 37 , S. Macfoy 28 , B. Machenschalk 10 , M. MacInnis 15 , D. M. Macleod 37 , I. Magaña Hernandez 21 , F. Magaña-Sandoval 45 , L. Magaña Zertuche 45 , R. M. Magee 64 , E. Majorana 36 , I. Maksimovic 138 , N. Man 67 , V. Mandic 46 , V. Mangano 47 , G. L. Mansell 26 , M. Manske 21,26 , M. Mantovani 31 , F. Marchesoni 44,53 , F. Marion 8 , S. Márka 51 , Z. Márka 51 , C. Markakis 12 , A. S. Markosyan 52 , A. Markowitz 1 , E. Maros 1 , A. Marquina 100 , F. Martelli 121,122 , L. Martellini 67 , I. W. Martin 47 , R. M. Martin 111 , D. V. Martynov 15 , K. Mason 15 , E. Massera 106 , A. Masserot 8 , T. J. Massinger 1 ,  \nM. Masso-Reid 47 , S. Mastrogiovanni 36,97 , A. Matas 46 , F. Matichard 1,15 , L. Matone 51 , N. Mavalvala 15 , N. Mazumder 69 , R. McCarthy 48 , D. E. McClelland 26 , S. McCormick 7 , L. McCuller 15 , S. C. McGuire 139 , G. McIntyre 1 , J. McIver 1 , D. J. McManus 26 , L. McNeill 6 , T. McRae 26 , S. T. McWilliams 42 , D. Meacher 64 , G. D. Meadors 10,39 , M. Mehmet 10 , J. Meidam 14 , E. Mejuto-Villa 9 , A. Melatos 96 , G. Mendell 48 , R. A. Mercer 21 , E. L. Merilh 48 , M. Merzougui 67 , S. Meshkov 1 , C. Messenger 47 , C. Messick 64 , R. Metzdorff 71 , P. M. Meyers 46 , H. Miao 59 , C. Michel 27 , H. Middleton 59 , E. E. Mikhailov 140 , L. Milano 4,79 , A. L. Miller 5,36,97 , B. B. Miller 89 , J. Miller 15 , M. Millhouse 101 , M. C. Milovich-Goff 109 , O. Minazzoli 67,141 , Y. Minenkov 34 , J. Ming 39 , C. Mishra 142 , S. Mitra 19 , V. P. Mitrofanov 62 , G. Mitselmakher 5 , R. Mittleman 15 , D. Moffa 85 , A. Moggi 24 , K. Mogushi 11 , M. Mohan 31 , S. R. P. Mohapatra 15 , M. Montani 121,122 , C. J. Moore 13 , D. Moraru 48 , G. Moreno 48 , S. R. Morriss 103 , B. Mours 8 , C. M. Mow-Lowry 59 , G. Mueller 5 , A. W. Muir 37 , Arunava Mukherjee 10 , D. Mukherjee 21 , S. Mukherjee 103 , N. Mukund 19 , A. Mullavey 7 , J. Munch 73 , E. A. Muñiz 45 , M. Muratore 38 , P. G. Murray 47 , K. Napier 77 , I. Nardecchia 33,34 , L. Naticchioni 36,97 , R. K. Nayak 143 , J. Neilson 109 , G. Nelemans 14,66 , T. J. N. Nelson 7 , M. Nery 10 , A. Neunzert 118 , L. Nevin 1 , J. M. Newport 124 , G. Newton 47, ‡ , K. K. Y. Ng 93 , T. T. Nguyen 26 , D. Nichols 66 , A. B. Nielsen 10 , S. Nissanke 14,66 , A. Nitz 10 , A. Noack 10 , F. Nocera 31 , D. Nolting 7 , C. North 37 , L. K. Nuttall 37 , J. Oberling 48 , G. D. O ' Dea 109 , G. H. Ogin 144 , J. J. Oh 131 , S. H. Oh 131 , F. Ohme 10 , M. A. Okada 16 , M. Oliver 102 , P. Oppermann 10 , Richard J. Oram 7 , B. O ' Reilly 7 , R. Ormiston 46 , L. F. Ortega 5 , R. O ' Shaughnessy 58 , S. Ossokine 39 , D. J. Ottaway 73 , H. Overmier 7 , B. J. Owen 84 , A. E. Pace 64 , J. Page 137 , M. A. Page 65 , A. Pai 116,145 , S. A. Pai 61 , J. R. Palamos 70 , O. Palashov 128 , C. Palomba 36 , A. Pal-Singh 35 , Howard Pan 87 , Huang-Wei Pan 87 , B. Pang 49 , P. T. H. Pang 93 , C. Pankow 89 , F. Pannarale 37 , B. C. Pant 61 , F. Paoletti 24 , A. Paoli 31 , M. A. Papa 10,21,39 , A. Parida 19 , W. Parker 7 , D. Pascucci 47 , A. Pasqualetti 31 , R. Passaquieti 23,24 , D. Passuello 24 , M. Patil 133 , B. Patricelli 24,146 , B. L. Pearlstone 47 , M. Pedraza 1 , R. Pedurand 27,147 , L. Pekowsky 45 , A. Pele 7 , S. Penn 148 , C. J. Perez 48 , A. Perreca 1,95,110 , L. M. Perri 89 , H. P. Pfeiffer 39,90 , M. Phelps 47 , O. J. Piccinni 36,97 , M. Pichot 67 , F. Piergiovanni 121,122 , V. Pierro 9 , G. Pillant 31 , L. Pinard 27 , I. M. Pinto 9 , M. Pirello 48 , M. Pitkin 47 , M. Poe 21 , R. Poggiani 23,24 , P. Popolizio 31 , E. K. Porter 40 , A. Post 10 , J. Powell 47,149 , J. Prasad 19 , J. W. W. Pratt 38 , G. Pratten 102 , V. Predoi 37 , T. Prestegard 21 , M. Prijatelj 10 , M. Principe 9 , S. Privitera 39 , G. A. Prodi 95,110 , L. G. Prokhorov 62 , O. Puncken 10 , M. Punturo 44 , P. Puppo 36 , M. Pürrer 39 , H. Qi 21 , V. Quetschke 103 , E. A. Quintero 1 , R. Quitzow-James 70 , F. J. Raab 48 , D. S. Rabeling 26 , H. Radkins 48 , P. Raffai 56 , S. Raja 61 , C. Rajan 61 , B. Rajbhandari 84 , M. Rakhmanov 103 , K. E. Ramirez 103 , A. Ramos-Buades 102 , P. Rapagnani 36,97 , V. Raymond 39 , M. Razzano 23,24 , J. Read 30 , T. Regimbau 67 , L. Rei 60 , S. Reid 63 , D. H. Reitze 1,5 , W. Ren 12 , S. D. Reyes 45 , F. Ricci 36,97 , P. M. Ricker 12 , S. Rieger 10 , K. Riles 118 , M. Rizzo 58 , N. A. Robertson 1,47 , R. Robie 47 , F. Robinet 29 , A. Rocchi 34 , L. Rolland 8 , J. G. Rollins 1 , V. J. Roma 70 , R. Romano 3,4 , C. L. Romel 48 , J. H. Romie 7 , D. Rosi ń ska 57,150 , M. P. Ross 151 , S. Rowan 47 , A. Rüdiger 10 , P. Ruggi 31 , G. Rutins 28 , K. Ryan 48 , S. Sachdev 1 , T. Sadecki 48 , L. Sadeghian 21 , M. Sakellariadou 152 , L. Salconi 31 , M. Saleem 116 , F. Salemi 10 , A. Samajdar 143 , L. Sammut 6 , L. M. Sampson 89 , E. J. Sanchez 1 , L. E. Sanchez 1 , N. Sanchis-Gual 25 , V. Sandberg 48 , J. R. Sanders 45 , B. Sassolas 27 , B. S. Sathyaprakash 37,64 , P. R. Saulson 45 , O. Sauter 118 , R. L. Savage 48 , A. Sawadsky 35 , P. Schale 70 , M. Scheel 49 , J. Scheuer 89 , J. Schmidt 10 , P. Schmidt 1,66 , R. Schnabel 35 , R. M. S. Scho /uniFB01 eld 70 , A. Schönbeck 35 , E. Schreiber 10 , D. Schuette 10,22 , B. W. Schulte 10 , B. F. Schutz 10,37 , S. G. Schwalbe 38 , J. Scott 47 , S. M. Scott 26 , E. Seidel 12 , D. Sellers 7 , A. S. Sengupta 153 , D. Sentenac 31 , V. Sequino 17,33,34 , A. Sergeev 128 , D. A. Shaddock 26 , T. J. Shaffer 48 , A. A. Shah 137 , M. S. Shahriar 89 , M. B. Shaner 109 , L. Shao 39 , B. Shapiro 52 , P. Shawhan 76 , A. Sheperd 21 , D. H. Shoemaker 15 , D. M. Shoemaker 77 , K. Siellez 77 , X. Siemens 21 , M. Sieniawska 57 , D. Sigg 48 , A. D. Silva 16 , L. P. Singer 80 , A. Singh 10,22,39 , A. Singhal 17,36 , A. M. Sintes 102 , B. J. J. Slagmolen 26 , B. Smith 7 , J. R. Smith 30 , R. J. E. Smith 1,6 , S. Somala 154 , E. J. Son 131 , J. A. Sonnenberg 21 , B. Sorazu 47 , F. Sorrentino 60 , T. Souradeep 19 , A. P. Spencer 47 , A. K. Srivastava 105 , K. Staats 38 , A. Staley 51 , M. Steinke 10 , J. Steinlechner 35,47 , S. Steinlechner 35 , D. Steinmeyer 10 , S. P. Stevenson 59,149 , R. Stone 103 , D. J. Stops 59 , K. A. Strain 47 , G. Stratta 121,122 , S. E. Strigin 62 , A. Strunk 48 , R. Sturani 155 , A. L. Stuver 7 , T. Z. Summerscales 156 , L. Sun 96 , S. Sunil 105 , J. Suresh 19 , P. J. Sutton 37 , B. L. Swinkels 31 , M. J. Szczepa ń czyk 38 , M. Tacca 14 , S. C. Tait 47 , C. Talbot 6 , D. Talukder 70 , D. B. Tanner 5 , M. Tápai 117 , A. Taracchini 39 , J. D. Tasson 72 , J. A. Taylor 137 , R. Taylor 1 , S. V. Tewari 148 , T. Theeg 10 , F. Thies 10 , E. G. Thomas 59 , M. Thomas 7 , P. Thomas 48 , K. A. Thorne 7 , K. S. Thorne 49 , E. Thrane 6 , S. Tiwari 17,95 , V. Tiwari 37 , K. V. Tokmakov 63 , K. Toland 47 , M. Tonelli 23,24 , Z. Tornasi 47 , A. Torres-Forné 25 , C. I. Torrie 1 , D. Töyrä 59 , F. Travasso 31,44 , G. Traylor 7 , J. Trinastic 5 , M. C. Tringali 95,110 , L. Trozzo 24,157 , K. W. Tsang 14 , M. Tse 15 , R. Tso 1 , L. Tsukada 82 , D. Tsuna 82 , D. Tuyenbayev 103 , K. Ueno 21 , D. Ugolini 158 , C. S. Unnikrishnan 119 , A. L. Urban 1 , S. A. Usman 37 , H. Vahlbruch 22 , G. Vajente 1 , G. Valdes 2 , N. van Bakel 14 , M. van Beuzekom 14 , J. F. J. van den Brand 14,75 , C. Van Den Broeck 14,159 , D. C. Vander-Hyde 45 , L. van der Schaaf 14 , J. V. van Heijningen 14 , A. A. van Veggel 47 , M. Vardaro 54,55 , V. Varma 49 , S. Vass 1 , M. Vasúth 50 , A. Vecchio 59 , G. Vedovato 55 , J. Veitch 47 , P. J. Veitch 73 , K. Venkateswara 151 , G. Venugopalan 1 , D. Verkindt 8 , F. Vetrano 121,122 , A. Viceré 121,122 , A. D. Viets 21 , S. Vinciguerra 59 , D. J. Vine 28 , J.-Y. Vinet 67 , S. Vitale 15 , T. Vo 45 , H. Vocca 43,44 , C. Vorvick 48 , S. P. Vyatchanin 62 , A. R. Wade 1 , L. E. Wade 85 , M. Wade 85 , R. Walet 14 , M. Walker 30 , L. Wallace 1 , S. Walsh 10,21,39 , G. Wang 17,122 , H. Wang 59 , J. Z. Wang 64 , W. H. Wang 103 , Y. F. Wang 93 , R. L. Ward 26 , J. Warner 48 , M. Was 8 , J. Watchi 98 , B. Weaver 48 , L.-W. Wei 10,22 , M. Weinert 10 , A. J. Weinstein 1 , R. Weiss 15 , L. Wen 65 , E. K. Wessel 12 , P. Weßels 10 , J. Westerweck 10 , T. Westphal 10 , K. Wette 26 , J. T. Whelan 58 , S. E. Whitcomb 1 , B. F. Whiting 5 ,  \n```\n, ,\n```  \n```\nC. Whittle 6 , D. Wilken 10 , D. Williams 47 , R. D. Williams 1 , A. R. Williamson 66 , J. L. Willis 1,160 , B. Willke 10,22 , M. H. Wimmer 10 , W. Winkler 10 , C. C. Wipf 1 , H. Wittel 10,22 , G. Woan 47 , J. Woehler 10 , J. Wofford 58 , K. W. K. Wong 93 , J. Worden 48 , J. L. Wright 47 , D. S. Wu 10 , D. M. Wysocki 58 , S. Xiao 1 , H. Yamamoto 1 , C. C. Yancey 76 , L. Yang 161 , M. J. Yap 26 M. Yazback 5 , Hang Yu 15 , Haocun Yu 15 , M. Yvert 8 , A. Zadro ż ny 132 , M. Zanolin 38 , T. Zelenova 31 , J.-P. Zendri 55 , M. Zevin 89 L. Zhang 1 , M. Zhang 140 , T. Zhang 47 , Y.-H. Zhang 58 , C. Zhao 65 , M. Zhou 89 , Z. Zhou 89 , S. J. Zhu 10,39 , X. J. Zhu 6 , A. B. Zimmerman 90 , M. E. Zucker 1,15 , J. Zweizig 1 , ( LIGO Scienti /uniFB01 c Collaboration and Virgo Collaboration ) E. Burns 80 , P. Veres 162 , D. Kocevski 137 , J. Racusin 80 , A. Goldstein 163 , V. Connaughton 163 , M. S. Briggs 162,164 , L. Blackburn 15,165 , R. Hamburg 162,164 , C. M. Hui 137 , A. von Kienlin 166 , J. McEnery 80 , R. D. Preece 162,164 , C. A. Wilson-Hodge 137 , E. Bissaldi 167,168 , W. H. Cleveland 163 , M. H. Gibby 169 , M. M. Giles 169 , R. M. Kippen 170 , S. McBreen 171 , C. A. Meegan 162 , W. S. Paciesas 163 , S. Poolakkil 162,164 , O. J. Roberts 163 , M. Stanbro 164 , ( Fermi Gamma-ray Burst Monitor ) and V. Savchenko 172 , C. Ferrigno 172 , E. Kuulkers 173 , A. Bazzano 174 , E. Bozzo 172 , S. Brandt 175 , J. Chenevez 175 , T. J.-L. Courvoisier 172 , R. Diehl 166 , A. Domingo 176 , L. Hanlon 177 , E. Jourdain 178 , P. Laurent 40,179 , F. Lebrun 180 , A. Lutovinov 181 , S. Mereghetti 182 , L. Natalucci 174 , J. Rodi 174 , J.-P. Roques 178 , R. Sunyaev 181,183 , and P. Ubertini 174 ( INTEGRAL ) † Deceased, 2017 February. ‡ Deceased, 2016 December. 1 LIGO, California Institute of Technology, Pasadena, CA 91125, USA 2 Louisiana State University, Baton Rouge, LA 70803, USA 3 Università di Salerno, Fisciano, I-84084 Salerno, Italy 4 INFN, Sezione di Napoli, Complesso Universitario di Monte S.Angelo, I-80126 Napoli, Italy 5 University of Florida, Gainesville, FL 32611, USA 6 OzGrav, School of Physics & Astronomy, Monash University, Clayton, VIC 3800, Australia 7 LIGO Livingston Observatory, Livingston, LA 70754, USA 8 Laboratoire d ' Annecy-le-Vieux de Physique des Particules ( LAPP ) , Université Savoie Mont Blanc, CNRS / IN2P3, F-74941 Annecy, France 9 University of Sannio at Benevento, I-82100 Benevento, Italy and INFN, Sezione di Napoli, I-80100 Napoli, Italy 10 Max Planck Institute for Gravitational Physics ( Albert Einstein Institute ) , D-30167 Hannover, Germany 11 The University of Mississippi, University, MS 38677, USA 12 NCSA, University of Illinois at Urbana-Champaign, Urbana, IL 61801, USA 13 University of Cambridge, Cambridge CB2 1TN, UK 14 Nikhef, Science Park, 1098 XG Amsterdam, The Netherlands 15 LIGO, Massachusetts Institute of Technology, Cambridge, MA 02139, USA 16 Instituto Nacional de Pesquisas Espaciais, 12227-010 São José dos Campos, São Paulo, Brazil 17 Gran Sasso Science Institute ( GSSI ) , I-67100 L ' Aquila, Italy 18 INFN, Laboratori Nazionali del Gran Sasso, I-67100 Assergi, Italy 19 Inter-University Centre for Astronomy and Astrophysics, Pune 411007, India 20 International Centre for Theoretical Sciences, Tata Institute of Fundamental Research, Bengaluru 560089, India 21 University of Wisconsin-Milwaukee, Milwaukee, WI 53201, USA 22 Leibniz Universität Hannover, D-30167 Hannover, Germany 23 Università di Pisa, I-56127 Pisa, Italy 24 INFN, Sezione di Pisa, I-56127 Pisa, Italy 25 Departamento de Astronomía y Astrofísica, Universitat de València, E-46100 Burjassot, València, Spain 26 OzGrav, Australian National University, Canberra, ACT 0200, Australia 27 Laboratoire des Matériaux Avancés ( LMA ) , CNRS / IN2P3, F-69622 Villeurbanne, France 28 SUPA, University of the West of Scotland, Paisley PA1 2BE, UK 29 LAL, Univ. Paris-Sud, CNRS / IN2P3, Université Paris-Saclay, F-91898 Orsay, France 30 California State University Fullerton, Fullerton, CA 92831, USA 31 European Gravitational Observatory ( EGO ) , I-56021 Cascina, Pisa, Italy 32 Chennai Mathematical Institute, Chennai 603103, India 33 Università di Roma Tor Vergata, I-00133 Roma, Italy 34 INFN, Sezione di Roma Tor Vergata, I-00133 Roma, Italy 35 Universität Hamburg, D-22761 Hamburg, Germany 36 INFN, Sezione di Roma, I-00185 Roma, Italy 37 Cardiff University, Cardiff CF24 3AA, UK 38 Embry-Riddle Aeronautical University, Prescott, AZ 86301, USA 39 Max Planck Institute for Gravitational Physics ( Albert Einstein Institute ) , D-14476 Potsdam-Golm, Germany 40 APC, AstroParticule et Cosmologie, Université Paris Diderot, CNRS / IN2P3, CEA / Irfu, Observatoire de Paris, Sorbonne Paris Cité, F-75205 Paris Cedex 13, France 41 Korea Institute of Science and Technology Information, Daejeon 34141, Korea 42 West Virginia University, Morgantown, WV 26506, USA 43\n```  \nUniversità di Perugia, I-06123 Perugia, Italy  \n- 44 INFN, Sezione di Perugia, I-06123 Perugia, Italy  \n45  \nSyracuse University, Syracuse, NY 13244, USA  \n- 46 University of Minnesota, Minneapolis, MN 55455, USA\n- 47 SUPA, University of Glasgow, Glasgow G12 8QQ, UK  \n48  \nLIGO Hanford Observatory, Richland, WA 99352, USA  \n49 Caltech CaRT, Pasadena, CA 91125, USA  \nWigner RCP, RMKI, H-1121 Budapest, Konkoly Thege Miklós út 29-33, Hungary  \n51 Columbia University, New York, NY 10027, USA  \n52 Stanford University, Stanford, CA 94305, USA  \n53 Università di Camerino, Dipartimento di Fisica, I-62032 Camerino, Italy  \nUniversità di Padova, Dipartimento di Fisica e Astronomia, I-35131 Padova, Italy  \n55  \nINFN, Sezione di Padova, I-35131 Padova, Italy  \n56 Institute of Physics, Eötvös University, Pázmány P. s. 1 / A, Budapest 1117, Hungary  \nNicolaus Copernicus Astronomical Center, Polish Academy of Sciences, 00-716, Warsaw, Poland  \n58 Rochester Institute of Technology, Rochester, NY 14623, USA  \n59  \nUniversity of Birmingham, Birmingham B15 2TT, UK  \n60  \nINFN, Sezione di Genova, I-16146 Genova, Italy  \n61  \nRRCAT, Indore MP 452013, India  \nFaculty of Physics, Lomonosov Moscow State University, Moscow 119991, Russia  \n63 SUPA, University of Strathclyde, Glasgow G1 1XQ, UK  \n- 64 The Pennsylvania State University, University Park, PA 16802, USA  \n65 OzGrav, University of Western Australia, Crawley, WA 6009, Australia  \n66 Department of Astrophysics / IMAPP, Radboud University Nijmegen, P.O. Box 9010, 6500 GL Nijmegen, The Netherlands  \n67  \nArtemis, Université Côte d ' Azur, Observatoire Côte d ' Azur, CNRS, CS 34229, F-06304 Nice Cedex 4, France 68  \nInstitut FOTON, CNRS, Université de Rennes 1, F-35042 Rennes, France  \n69 Washington State University, Pullman, WA 99164, USA 70  \nUniversity of Oregon, Eugene, OR 97403, USA  \n71 Laboratoire Kastler Brossel, UPMC-Sorbonne Universités, CNRS, ENS-PSL Research University, Collège de France, F-75005 Paris, France 72  \n73  \nCarleton College, North /uniFB01 eld, MN 55057, USA  \nOzGrav, University of Adelaide, Adelaide, SA 5005, Australia  \nAstronomical Observatory Warsaw University, 00-478 Warsaw, Poland  \n74  \n75 VU University Amsterdam, 1081 HV Amsterdam, The Netherlands  \n76 University of Maryland, College Park, MD 20742, USA  \n77 Center for Relativistic Astrophysics, Georgia Institute of Technology, Atlanta, GA 30332, USA  \n78  \nUniversité Claude Bernard Lyon 1, F-69622 Villeurbanne, France  \n79 Università di Napoli ' Federico II, ' Complesso Universitario di Monte S.Angelo, I-80126 Napoli, Italy  \n80  \nNASA Goddard Space Flight Center, Greenbelt, MD 20771, USA  \n81 Dipartimento di Fisica, Università degli Studi di Genova, I-16146 Genova, Italy  \n82 RESCEU, University of Tokyo, Tokyo, 113-0033, Japan 83  \n84  \nTsinghua University, Beijing 100084, China  \nTexas Tech University, Lubbock, TX 79409, USA  \n85 Kenyon College, Gambier, OH 43022, USA  \n86 Museo Storico della Fisica e Centro Studi e Ricerche Enrico Fermi, I-00184 Roma, Italy 87  \nNational Tsing Hua University, Hsinchu City, 30013 Taiwan, Republic of China  \n88 Charles Sturt University, Wagga Wagga, NSW 2678, Australia  \nCenter for Interdisciplinary Exploration & Research in Astrophysics  \n90  \n107  \n109  \n112  \n(  \nCIERA  \n)  \n, Northwestern University, Evanston, IL 60208, USA  \nCanadian Institute for Theoretical Astrophysics, University of Toronto, Toronto, ON M5S 3H8, Canada  \n91 University of Chicago, Chicago, IL 60637, USA  \n92 Pusan National University, Busan 46241, Korea  \n93 The Chinese University of Hong Kong, Shatin, NT, Hong Kong  \n94  \nINAF, Osservatorio Astronomico di Padova, I-35122 Padova, Italy 95 INFN, Trento Institute for Fundamental Physics and Applications, I-38123 Povo, Trento, Italy 96  \nOzGrav, University of Melbourne, Parkville, VIC 3010, Australia  \n- 97 Università di Roma ' La Sapienza, ' I-00185 Roma, Italy  \n98  \n99  \nUniversité Libre de Bruxelles, Brussels B-1050, Belgium  \nSonoma State University, Rohnert Park, CA 94928, USA  \nDepartamento de Matemáticas, Universitat de València, E-46100 Burjassot, València, Spain  \n101 Montana State University, Bozeman, MT 59717, USA  \nUniversitat de les Illes Balears, IAC3  \n-  \nIEEC, E-07122 Palma de Mallorca, Spain  \n103  \nThe University of Texas Rio Grande Valley, Brownsville, TX 78520, USA 104  \n105  \nBellevue College, Bellevue, WA 98007, USA  \nInstitute for Plasma Research, Bhat, Gandhinagar 382428, India  \n106 The University of Shef /uniFB01 eld, Shef /uniFB01 eld S10 2TN, UK  \nDipartimento di Scienze Matematiche, Fisiche e Informatiche, Università di Parma, I-43124 Parma, Italy  \n108 INFN, Sezione di Milano Bicocca, Gruppo Collegato di Parma, I-43124 Parma, Italy  \nCalifornia State University, Los Angeles, 5151 State University Dr, Los Angeles, CA 90032, USA  \n110  \nUniversità di Trento, Dipartimento di Fisica, I-38123 Povo, Trento, Italy 111  \nMontclair State University, Montclair, NJ 07043, USA  \nNational Astronomical Observatory of Japan, 2-21-1 Osawa, Mitaka, Tokyo 181-8588, Japan  \n113  \nObservatori Astronòmic, Universitat de València, E-46980 Paterna, València, Spain  \n114  \nSchool of Mathematics, University of Edinburgh, Edinburgh EH9 3FD, UK  \nUniversity and Institute of Advanced Research, Koba Institutional Area, Gandhinagar Gujarat 382007, India  \n116  \nIISER-TVM, CET Campus, Trivandrum Kerala 695016, India  \n117  \nUniversity of Szeged, Dóm tér 9, Szeged 6720, Hungary  \n118 University of Michigan, Ann Arbor, MI 48109, USA  \n89  \n115  \n57  \n100  \n50  \n54  \n62  \n102  \n- 119 Tata Institute of Fundamental Research, Mumbai 400005, India  \nINAF, Osservatorio Astronomico di Capodimonte, I-80131 Napoli, Italy  \n120  \n121  \n122  \nUniversità degli Studi di Urbino  \n'  \nCarlo Bo,  \n'  \nI-61029 Urbino, Italy  \nINFN, Sezione di Firenze, I-50019 Sesto Fiorentino, Firenze, Italy  \nPhysik-Institut, University of Zurich, Winterthurerstrasse 190, 8057 Zurich, Switzerland  \n124  \nAmerican University, Washington, DC 20016, USA  \n125  \nUniversity of Bia  \nł  \nystok, 15-424 Bia  \nł  \nystok, Poland  \nUniversity of Southampton, Southampton SO17 1BJ, UK  \n126  \nUniversity of Washington Bothell, 18115 Campus Way NE, Bothell, WA 98011, USA  \n128  \n129  \n131  \nInstitute of Applied Physics, Nizhny Novgorod, 603950, Russia  \nKorea Astronomy and Space Science Institute, Daejeon 34055, Korea  \n130  \nInje University Gimhae, South Gyeongsang 50834, Korea  \nNational Institute for Mathematical Sciences, Daejeon 34047, Korea  \n132  \nNCBJ, 05-400  \nŚ  \nwierk-Otwock, Poland  \nInstitute of Mathematics, Polish Academy of Sciences, 00656 Warsaw, Poland  \n134  \nHillsdale College, Hillsdale, MI 49242, USA  \n135  \n136  \nHanyang University, Seoul 04763, Korea  \nSeoul National University, Seoul 08826, Korea  \nNASA Marshall Space Flight Center, Huntsville, AL 35812, USA  \n138  \nESPCI, CNRS, F-75005 Paris, France  \nSouthern University and A&M College, Baton Rouge, LA 70813, USA  \n140  \nCollege of William and Mary, Williamsburg, VA 23187, USA  \nCentre Scienti  \n142  \n/uniFB01  \nque de Monaco, 8 quai Antoine Ier, MC-98000, Monaco  \nIndian Institute of Technology Madras, Chennai 600036, India  \n143  \nIISER-Kolkata, Mohanpur, West Bengal 741252, India  \nWhitman College, 345 Boyer Avenue, Walla Walla, WA 99362, USA  \n137  \n139  \n141  \n144  \nIndian Institute of Technology Bombay, Powai, Mumbai, Maharashtra 400076, India  \n146  \nScuola Normale Superiore, Piazza dei Cavalieri 7, I-56126 Pisa, Italy  \n147  \nUniversité de Lyon, F-69361 Lyon, France  \nHobart and William Smith Colleges, Geneva, NY 14456, USA  \n148  \nOzGrav, Swinburne University of Technology, Hawthorn VIC 3122, Australia  \n149  \nJanusz Gil Institute of Astronomy, University of Zielona Góra, 65-265 Zielona Góra, Poland  \n151  \nUniversity of Washington, Seattle, WA 98195, USA  \ns College London, University of London, London WC2R 2LS, UK  \n152  \nKing  \n'  \nIndian Institute of Technology, Gandhinagar Ahmedabad Gujarat 382424, India  \n153  \n154 Indian Institute of Technology Hyderabad, Sangareddy, Khandi, Telangana 502285, India  \nInternational Institute of Physics, Universidade Federal do Rio Grande do Norte, Natal RN 59078-970, Brazil  \n156  \n158  \nAndrews University, Berrien Springs, MI 49104, USA  \n157  \nUniversità di Siena, I-53100 Siena, Italy  \nTrinity University, San Antonio, TX 78212, USA  \nVan Swinderen Institute for Particle Physics and Gravity, University of Groningen, Nijenborgh 4, 9747 AG Groningen, The Netherlands  \n160 Abilene Christian University, Abilene, TX 79699, USA  \n161  \nColorado State University, Fort Collins, CO 80523, USA  \nCenter for Space Plasma and Aeronomic Research, University of Alabama in Huntsville, 320 Sparkman Drive, Huntsville, AL 35899, USA  \n163  \nScience and Technology Institute, Universities Space Research Association, Huntsville, AL 35805, USA  \nSpace Science Department, University of Alabama in Huntsville, 320 Sparkman Drive, Huntsville, AL 35899, USA  \n165  \nHarvard-Smithsonian Center for Astrophysics, 60 Garden St, Cambridge, MA 02138, USA  \nMax-Planck-Institut für extraterrestrische Physik, Giessenbachstrasse 1, D-85748 Garching, Germany  \n167  \n168  \nPolitecnico di Bari, Via Edoardo Orabona, 4, I-70126 Bari BA, Italy  \nIstituto Nazionale di Fisica Nucleare, Sezione di Bari, I-70126 Bari, Italy  \n169  \nJacobs Technology, Inc., Huntsville, AL 35806, USA  \n- 170 Los Alamos National Laboratory, PO Box 1663, Los Alamos, NM 87545, USA  \n171 School of Physics, University College Dublin, Bel /uniFB01 eld, Stillorgan Road, Dublin 4, Ireland  \nISDC, Department of Astronomy, University of Geneva, Chemin d  \n'  \nÉcogia, 16 CH-1290 Versoix, Switzerland  \n164  \n172  \nEuropean Space Research and Technology Centre  \n(  \nESA  \n/  \nESTEC  \n)  \n, Keplerlaan 1, 2201 AZ Noordwijk, The Netherlands  \n- 174 INAF-Institute for Space Astrophysics and Planetology, Via Fosso del Cavaliere 100, I-00133 Roma, Italy 175  \n173  \nDTU Space -National Space Institute Elektrovej -Building 327 DK-2800 Kongens Lyngby Denmark 176 Centro de Astrobiología ( CAB-CSIC / INTA, ESAC Campus ) , Camino bajo del Castillo S / N, E-28692 Villanueva de la Cañada, Madrid, Spain  \n177  \n178  \nIrfu  \n179  \nDSM  \n/  \n/  \nSpace Science Group, School of Physics, University College Dublin, Bel  \n/uniFB01  \neld, Dublin 4, Ireland  \nIRAP, Université de Toulouse; CNRS; UPS; CNES; 9 Av. Roche, F-31028 Toulouse, France  \nService d  \n'  \nAstrophysique, Bat. 709 Orme des Merisiers CEA Saclay, F-91191 Gif-sur-Yvette Cedex, France  \nAPC, AstroParticule et Cosmologie, Université Paris Diderot, CNRS  \n181  \n/  \nIN2P3, CEA  \n/  \nIrfu, Observatoire de Paris Sorbonne Paris Cité, France  \nSpace Research Institute of Russian Academy of Sciences, Profsoyuznaya 84  \n182  \n/  \n32, 117997 Moscow, Russia  \nINAF, IASF-Milano, via E.Bassini 15, I-20133 Milano, Italy  \n183 Max Planck Institute for Astrophysics, Karl-Schwarzschild-Str. 1, Garching b. Munchen D-85741, Germany  \n159  \n162  \n180  \n155  \n166  \n123  \n127  \n133  \n145  \n150"}