The publicly available Chandra observations are described in Section 2 . In section Section 3 we discuss the interpretation of the X-ray data. The main conclusions are summarized in Section 4 .

Of considerable interest is the nature of the compact remnant left behind by this event. Some theoretical works have suggested that the compact remnant is unlikely to be a neutron star that lived longer than a few months (Margalit & Metzger 2017 ; Villar et al. 2017 ). Others do allow a neutron star remnant with a relatively modest magnetic field to be consistent with early X-ray measurements (Yu & Dai 2017 , but see Sections 3.2 and 4 ). Observations published to date are inconclusive. Here, we address anew the final product of the merger, and if and how observations could rule out or confirm either a neutron star or black hole compact remnant. We also discuss the unresolved issue of the nature of the propagating external shock driven by the merger. We requested Director's Discretionary Time Chandra observations in an attempt to address these issues.

The discovery of gravitational waves (GWs) from a binary neutron star merger by the Laser Interferometer Gravitational-wave Observatory (LIGO; Abbott et al. 2017a ), the associated electromagnetic signal that arrived with a delay of 1.74 s (Abbott et al. 2017b ), and the precise localization (Coulter et al. 2017 ) have opened a new and exciting frontier that should provide answers to long-standing questions regarding the synthesis of r-process elements (e.g., Kasen et al. 2017 ; Tanvir et al. 2017 ) and the generation of relativistic jets and γ-ray photons (Kasliwal et al. 2017 ) in these events and possibly other high-energy sources. X-ray observations with the Chandra X-ray Observatory initially did not detect the event at two to three days post merger, but subsequent observations showed a brightening X-ray source nine days after the merger (Haggard et al. 2017 ; Margutti et al. 2017 ; Troja et al. 2017 ).

Given the poor fit with a single power law and the suggestive evidence for fading from the K–S test, we next fit a broken power-law model, before t break and after t break . Because there is little data after a possible break to constrain the decay slope, we fix a 2 to one of two values: −0.9 if the shock front is moving at relativistic speeds and −1.2 if the shock front is sub-relativistic. For a 2 = −0.9, we find with days. For , we find with days. The two fits have χ 2 values of 5.6 and 5.4, respectively, for 6 degrees of freedom. We plot these best fits in Figure 2 .

Figure 3. Cumulative distribution of the photon arrival times from the January observations concatenated together, with faint gray lines marking the divisions between observations. This is compared to a uniform distribution with a K–S test.

The likely reason for the poorness of these fits is the possible fading trend seen in the January observations, as suggested by D'Avanzo et al. ( 2018 ). In the most general case, we nonparametrically test whether the January observations are consistent with a constant count rate using a K–S test. The photon arrival times (during "good-time intervals") were compared to a uniform distribution. The five individual observations were concatenated into one, with the next observation beginning where the previous left off. The K–S probability that the photon arrival times are drawn from a uniform distribution is only 2.1%. Figure 3 shows the cumulative distributions with faint gray lines marking the observation divisions. The maximum difference occurs near the division between the third and fourth January observations.

Figure 2. Chandra light curve of the unabsorbed 0.5–8 keV flux from the available observations of GW170817. Uncertainties in flux take account of uncertainties in the spectral fits; see the text for details. The hatched region indicates when GW170817 was behind the Sun and unobservable with Chandra. The green line indicates the best-fit single power-law to the data, with a 1 = 0.5. The red and blue lines show the best-fit broken power laws, before t break and after. The red line shows the best fit when a 2 is fixed at −0.9 (expected when there is no energy being added to the forward shock and the shock front is moving at a relativistic speed), and the blue line shows the best fit when a 2 is fixed at −1.2 (expected for a sub-relativistic shock); in both cases, the best-fit a 1 is 0.7.

Although there is a considerable amount of stochastic variation in the calculated fluxes (and count-rates), there does seem to be an indication that the temporal behavior at days 153–163 is qualitatively different than at earlier times. If we take the two December observations together as a single flux point and the five January observations together as a single flux point, we fit the four-point X-ray light curve with a single power-law model in time, , and find that a 1 = 0.5 with a χ 2 of 7.8 for 2 degrees of freedom. If we use all of the flux measurements separately in the fit, we find a 1 = 0.5 with a χ 2 of 15.6 for 7 degrees of freedom. Neither of these fits is acceptable (p-values of 0.02 and 0.03, respectively), but we plot a power law with index 0.5 as the green line in Figure 2 for reference.

The large uncertainties on the power-law indices make it difficult to determine whether or not the X-ray spectrum has changed slope. To address this, we perform pairwise Kolmogorov–Smirnov (K–S) tests on the detected photon energies and find marginal evidence for spectral evolution. The K–S tests show probabilities of the detected photon energies being drawn from the same parent distribution in the ~10%–100% range, which indicates consistency of the observed spectra. The smallest probability for being drawn from the same parent distribution is for the first and fourth observations (ObsIDs 19294 and 20861), which have a 6% probability of being drawn from the same parent distribution; this is not convincingly small enough to claim a spectral change.

We also perform the fits with no additional absorption component. Compared to the fits allowing for additional absorption, these fits have a slightly worse fit statistic, slightly lower best-fit power-law indices, and roughly the same unabsorbed fluxes (Table 2 ).

At the midpoint of the observation(s) being fit. For the top set of fits, the best-fit additional column isand, for the bottom set of fits, isPL index is the parameter β, where the number counts spectrum,

Notes. The fits on the left allow for additional absorption beyond that through the Milky Way, while the fits on the right do not. The top sets of fits allow different power-law normalizations and slopes for each Chandra observation listed in Table 1 . The bottom sets of fits constrain the two observations of the December epoch (days 108 and 111) to have the same slope and normalization and the five observations of the January epoch (between days 153 and 164) to have the same slope and normalization.

In our fits, we require the same absorbing column for all spectra but allow the nine power-law indices and power-law normalizations to vary independently. In addition, we perform another fit with both of the December observations required to have the same power-law index and normalization and all of the January observations required to have the same power-law index and normalization. This resulted in a slightly different best-fit value for the additional column density. The results are given in Table 2 , separated by a horizontal rule for the two sets of simultaneous fits. All uncertainties are 1σ-equivalent (68%) confidence intervals. The reported fluxes are integrated from the unabsorbed models. Uncertainties on those fluxes are calculated as the 68%-confidence bounds of the integrated, unabsorbed fluxes from Monte Carlo realizations (1000 samples) of the best-fit models, taking into account the uncertainties in the best-fit parameters (using the sample_flux command in Sherpa).

We performed a simultaneous fit of all nine unbinned source spectra in the 0.5–8 keV band with Sherpa using the modified Cash ( 1979 ) statistic cstat and the simplex optimization method. We fit the data with a power-law model with two absorption components. We use the Tuebingen-Boulder interstellar medium (ISM) absorption model (Wilms et al. 2000 ) and fix one absorbing column to the Galactic value of n H = 7.20 × 10 20 cm −2 calculated from the Effelsberg-Bonn H i Survey (Winkel et al. 2016 ) using the online tool at the Argelander-Institut für Astronomie. 6 We let the column density of the other absorption component vary to allow for absorption in the host galaxy and any absorption local to the event.

In each observation, we extracted spectra from a 1 0 radius region centered on GW170817 and a 1' source-free background region to the northwest. The net counts and count rate are given in Table 1 .

Note. Horizontal rules separate epochs. The last lines in the December and January epochs give the relevant information when considering all of the observations from the epoch as a single data point.

Figure 1. Chandra images of the field of GW170817. Each image is 30'' × 30'' and made with counts in the range 0.5–8 keV. The left image shows the combined data from August 26 and September 1 data (presented in Troja et al. 2017 , and reanalyzed and discussed here). The middle image shows the combined data from December 3 and 6 presented and discussed here. The right image shows the combined data from January 17–28 presented and discussed here. The colorbar indicates the number of counts in each pixel.

We list in Table 1 the starting dates of all publicly available observations, their ObsIDs, and their exposure times in ks. All observations were taken with the telescope aimpoint on the Advanced CCD Imaging Spectrometer (ACIS) S3 chip. Data reduction was performed with the chandra_repro script, part of the Chandra Interactive Analysis of Observations (CIAO) software. We used CIAO version 4.9 and calibration database (CALDB) version 4.7.6. Images of the field of GW170817 are shown in Figure 1 .

In Section 3.1 we conclude, as have others (Haggard et al. 2017; Margutti et al. 2017; Troja et al. 2017; Mooley et al. 2018; Ruan et al. 2018), that the X-ray data from about 9 to 111 days after the GW detection are consistent with an expanding shocked plasma emitting synchrotron radiation due to its interaction with the ISM, and with energy input to the plasma during this time. The possible fading seen in Chandra data from days 153–163 would be consistent with reduced or no energy input. In Section 3.2, we show that considerations of X-ray optical depth give the time after which emission from a recently formed magnetic neutron star or a pulsar wind nebula (PWN) would dominate. The X-ray flux measured during days 107–111 and 153–163 is much smaller than that expected if a neutron star remnant were formed and lived for 102 days. This suggests that a black hole was born in this merger well before the observations on day 107. We propose other stringent tests of this conclusion in Section 4.

3.1. External Shock Origin of X-Ray Light Curve

The radio spectrum reported in Mooley et al. (2018) gave a relatively precise log slope of −0.61 ± 0.05, i.e., f ν ∝ ν−0.61±0.05, for data through day 107. Our best-fit X-ray power-law index, though poorly constrained, is consistent within uncertainties with those in the radio. Margutti et al. (2018) extended the single consistent power law conclusion to day 160. The most conservative scenario, therefore, is that the radio and the X-ray photons are produced in the same source via the synchrotron radiation mechanism. The synchrotron cooling frequency in a relativistic shock wave (e.g., Panaitescu & Kumar 2001; Granot & Sari 2002) can be shown to be where E is the energy in the shocked plasma, n is the number density of particles in the medium in the vicinity of the neutron star merger, B is the fraction of the thermal energy of the shocked plasma in magnetic fields, t is the time since the merger event in the observer frame, and Y is the Compton Y-parameter, which is less than 1 based on radio and X-ray data; integer subscript for a variable X, i.e., X m , is a convenient shorthand notation for X/10m in cgs units. The density n is expected to be of order 0.1 cm−3 or less in the ISM of the host galaxy of GW170817. Thus, even a year after the merger. Our X-ray observations are below the expected cooling frequency. The observed synchrotron flux at 1 keV is given by (e.g., Kumar & Zhang 2015) where e is the fraction of shock energy given to electrons. In deriving the above expression we have taken the electron energy distribution to be with p = 2.2 as suggested by the X-ray and radio spectra that give . For constant shock energy, Equation (2) suggests that the X-ray flux should decline with time. Instead, the X-ray and radio light curves are rising with time from day nine until at least day 111. For , we find from Equation (2) that energy in the shocked matter is increasing with time as The flux at 1 keV at 102 days (1.4 nJy) is larger than the flux at nine days by a factor of ~5, implying that α ~ 0.7. We thus conclude from Equation (3) that the energy in the shocked plasma increased by a factor ~20 between these two epochs. There are two ways that the energy in the external shock could increase with time. One is that the compact remnant left over in the merger continues to pump energy into the outward moving shock front. The other is that energy is supplied to the part of the external shock wave that we see by slower-moving ejecta or a structured jet.7 The first possibility, continuous production of energy from the compact object, is unlikely to work. This is because there is ~3 × 10−2 of debris between the central compact remnant and the external shock front (e.g., Kasliwal et al. 2017). If the central compact object produces ~1052 erg of energy, and a good fraction of this energy is used to accelerate the debris, then the speed of the debris would be ~0.7c. This speed is smaller than the outward velocity of the relativistic shock front. Any such energy will not be added to the shock until the shock speed has finally fallen below ~0.7c. This is expected to take a year or more (Section 3.2). Observations carried out then would be able to constrain the total energy in the outflow, as discussed in Section 4. We are thus left with the second possibility: at 107 days, there is much more energy in the shocked medium within our cone of visibility, i.e., , than there was at nine days, where γ is the Lorentz factor of the shocked fluid within . There are two possible ways this could be true: slower-moving ejecta caught up with the external shock, or more of the shock is visible as it decelerates and beaming effects decrease. We first consider the scenario where the unexpected shock energy is supplied by slower-moving, perhaps quasi-spherical, ejecta that catches up with the decelerating shock front. Let us take the distribution of energy with ejecta velocity (βγ) such that where the exponent is determined from the observed light curves. Combining the energy equation (e.g., Granot & Sari 2002) with the relation between the shock front radius r and observer time t yields Using Equation (4) we obtain which is consistent with Equation (3) provided that The X-ray data give α ~ 0.7, and therefore the rising light curve can be explained if , which is roughly consistent with the conclusion of Mooley et al. (2018). A rising light curve can also occur when E is a rapidly decreasing function of angle θ measured from the jet axis and the observer is located off axis, as was pointed out by Lazzati et al. (2017) for GW170817. In the case of either a top-hat structure or a more gently shaped angular structure, an off-axis observer will see a rising light curve as the external shock decelerates and the Lorentz beaming decreases. The rate of increase for a top-hat jet scales as t3, which is inconsistent with the data (Mooley et al. 2018; Ruan et al. 2018). For a structured jet, the rate of rise depends on how the energy declines with angle. In this case, an off-axis observer finds the energy in the shocked medium to increase with time because photons are detected from a larger fraction of the external shock. It can be shown that the rising X-ray flux between ~10 and 102 days ( ) can arise from a structured jet when , (θ c is the angular size of the jet core) with θ c ~ 3°, and the observer is located off axis at an angle ~30° with respect to the jet axis. The data for GW170817 are consistent with both of these possibilities: slower-moving ejecta that caught up with the external shock, or beaming effects from a structured jet. This degeneracy may be removed by careful modeling of afterglow light curve data about a year after the merger. If the rise is due to a structured jet with angular size of the core ~3°, then the X-flux should stop increasing at t ~ 102 days; there is evidence for this flattening in the Chandra data at around day 160 (Figure 2). We note that the total energy in the shocked ISM is highly uncertain because of the unknown parameters n, B , and e . We return to the discussion of the nature of the external shock in Section 4.

3.2. Merged Remnant a Black Hole or a Neutron Star?