X-RAY AND ULTRAVIOLET LINE EMISSION FROM SNR 1987A

Transcription

1 THE ASTROPHYSICAL JOURNAL, 476 : L31 L34, 1997 February The American Astronomical Society. All rights reserved. Printed in U.S.A. X-RAY AND ULTRAVIOLET LINE EMISSION FROM SNR 1987A KAZIMIERZ J. BORKOWSKI AND JOHN M. BLONDIN Department of Physics, North Carolina State University, Raleigh, NC AND RICHARD MCCRAY JILA, University of Colorado, Boulder, CO Received 1996 August 20; accepted 1996 November 27 ABSTRACT The soft X-ray emission seen from SN 1987A and the apparent deceleration of the radio source expansion suggest that the supernova blast wave has encountered a moderately dense H II region interior to the inner circumstellar ring. We simulate the hydrodynamics of this interaction and calculate the resulting X-ray and ultraviolet emission-line spectrum and light curves. The soft X-ray spectrum is dominated by emission lines of hydrogenic and helium-like C, N, O, and Ne; it is consistent with the ROSAT observations if Fe is depleted on grains. N V 1240 emission should be observable easily with the Hubble Space Telescope. The blast wave should strike the inner circumstellar ring around A.D Subject headings: circumstellar matter hydrodynamics supernovae: individual (SN 1987A) ultraviolet: ISM X-rays: ISM 1. INTRODUCTION The debris of SN 1987A is driving a blast wave into its circumstellar environment. The radio emission seen by the Australia Telescope (Gaensler et al. 1997) probably comes from relativistic electrons accelerated in the region between the blast wave and the reverse shock that is propagating into the supernova debris (Chevalier 1992). If so, the radio images suggest that the blast wave decelerated when it encountered relatively dense gas beginning at t days after outburst and that the interaction region had a toroidal rather than a spherical symmetry. Moreover, soft X-rays from SN 1987A were first detected by ROSAT at t days (Beuermann, Brandt, & Pietsch 1994; Gorenstein, Hughes, & Tucker 1994) and have continued to brighten since then (Hasinger, Aschenbach, & Trümper 1996). Chevalier & Dwarkadas (1995) have suggested that the radio and X-ray observations can be explained with a model in which the blast wave entered a relatively dense (n cm 3 ) H II region that separates the shocked stellar wind of the supernova progenitor from the inner circumstellar ring. They estimate that the blast wave will strike this ring around A.D H 2. When it does, the inner ring will become a bright source of optical, ultraviolet, and X-ray emission lines (Luo, McCray, & Slavin 1994; Borkowski, Blondin, & McCray 1997). But, in the meantime, SN 1987A is a steadily brightening source of X-rays. What can we infer from the current observations? What will we learn from future observations with AXAF and ASTRO-E? Can the present shock interaction be observed with the Hubble Space Telescope? As the reader will shortly see, the answers to these questions are very exciting indeed. 2. HYDRODYNAMICS We consider a model in which the supernova envelope drives a blast wave into its circumstellar environment. We L31 assume that the outer envelope of the supernova has a density profile given by 1 t v 10 yr 3 1 9, (1) 10 4 km s where the coefficient 1 is determined by a fit to the early photospheric light curve and spectrum (Eastman & Kirshner 1989) and has values ranging from amu cm 3 (model 14E1 of Shigeyama & Nomoto 1990) to amu cm 3 (model 10H of Woosley 1988), the value that we adopt here. Observations from the Australia Telescope (Gaensler et al. 1997) show that nonthermal radio emission from SN 1987A reappeared at t days and has continued to brighten since then. The radio images are consistent with an expanding toroidal source in the equatorial plane of the ring. The apparent expansion has decelerated, from average velocity V 1 30,000 km s 1 for t 1200 days to V H 400 km s 1 for 1200 days t 3200 days. At t 3200 days, the brightness peaks of the radio emission were located at radial distance r 2 0.8R r, where R r is the radius of the inner ring. To model the radio and X-ray observations, we have simulated the hydrodynamics of the impact of a supernova envelope (eq. [1]) with circumstellar gas consisting of three components. The first component is the undisturbed stellar wind of the blue supergiant progenitor, with a ratio of mass-loss rate to terminal velocity Ṁ/v w 10 7 M J yr 1 /(300 km s 1 ). The stellar wind terminates at a reverse wind shock at a radius cm, beyond which lies the second component, a uniform density bubble of shocked stellar wind with density n amu cm 3. The third component is a thick circular torus, with uniform density n II 150 amu cm 3, representing the H II region (cf. Chevalier & Dwarkadas 1995). The torus has a major radius R M R R cm and minor radius R m 0.5R M, so that the inner boundary is at a radius R II 0.5R M (the inner circumstellar ring is embedded within this torus). As the envelope propagates through the stellar wind, the

2 L32 BORKOWSKI, BLONDIN, & MCCRAY Vol. 476 FIG. 2. Emission measure (solid curves) and mass flux (dashed curves) vs. time. Contributions from the blast wave (upper solid and dashed curves) dominate those from the reverse shock (lower solid and dashed curves). interaction will establish a spherical blast wave and reverse shock conforming to Chevalier s (1982) similarity solution with s 2 and n 9. At t 950 days, the blast wave (B) encounters the reverse wind shock (R). Shortly thereafter, at t days, the blast wave enters the H II region and slows down suddenly as it crushes the inner boundary of the H II region. The mean velocity of the blast wave in the equatorial plane for 1200 days = t = 3200 days is V 4100 km s 1, somewhat greater than the value observed by Gaensler et al. (1997). Figure 1 (Plate L2) is a snapshot of a moment (t 9 yr) in our two-dimensional (azimuthally symmetric) hydrodynamical simulation of this interaction. One sees that the forward shock has reached roughly halfway around the exterior of the torus, while the transmitted shock has propagated through about 30% of the radius of the torus. The instability of the contact discontinuity separating the shocked circumstellar matter and the shocked supernova ejecta (cf. Chevalier, Blondin, & Emmering 1992) is evident. In this simulation, the blast wave will first strike the dense inner ring at t 20.5 yr, or A.D. 2007, at which time a very bright display of infrared, optical, ultraviolet, and X-ray emission will ensue (Luo et al. 1994; Borkowski et al. 1997). Figure 2 shows the evolution of the mass fluxes, Ṁ B (t) and Ṁ R (t), and emission measures, EM B (t) and EM R (t), of the shocked circumstellar gas and supernova envelope, respectively. Note that Ṁ B (t) and EM B (t) begin to increase earlier than Ṁ R (t) and EM R (t), when the blast wave enters the H II region, but Ṁ R (t) and EM R (t) begin to catch up at t 2 5 yr, when shocks reflected from the H II region merge with the original reverse shock and cause it to accelerate inward. 3. X-RAY EMISSION The X-ray spectrum of SN 1987A in the ROSAT band is soft; it must be dominated by line emission. Fitting the ROSAT position sensitive proportional counter (PSPC) pulse-height spectrum with a thermal plasma model (Raymond & Smith 1977) with cosmic abundances Z 0.3 and absorption by interstellar gas with solar abundances and column density N H cm 2, Hasinger et al. (1996) infer a plasma temperature kt 0.99 H 0.17 kev and emission measure EM (1.4 H 0.4) cm 3 (at t days). The fit is 2 not good, however ( red 2 3). FIG. 3. X-ray spectrum for a 5000 km s 1 shock, 3 yr after the impact with the H II region, corrected for interstellar absorption with N H cm 2. The spectrum is dominated by strong Ly and He lines of H-like and He-like N, O, and Ne ions. Are the ROSAT observations consistent with the Chevalier & Dwarkadas (1995) model? To test this hypothesis, we calculated the X-ray emission from a fast (v S km s 1 ) steady shock. (The results are fairly insensitive to velocity.) We assume that electrons are heated by Coulomb collisions with ions, without any collisionless heating at the shock front. The electron temperature, T e, is then much less than the ion temperature, T i, and is given approximately by Hamilton & Sarazin (1984): T e T 3/2 2ln 2/5 nt, (2) 503 T where ln 2 30 is the Coulomb logarithm, 3kT/2 is the mean energy per particle, and n is the total particle density. Equation (2) gives n kt e 3.8 kev e t v cm s 2/5 1 4/5 S. (3) km s The soft X-ray line emission seen by ROSAT comes primarily from gas with n e t cm 3 s, or kt e 1 2 kev, while ultraviolet resonance line emission comes primarily from gas with n e t cm 3 s, or kt e kev (see below). A Raymond-Smith (1977) model is not valid for such a flow ionization is far from coronal equilibrium but we find nevertheless that the emission measures required to account for the observed soft X-ray emission are approximately equal to those quoted by Hasinger et al. (1996). To estimate the light curves of emission lines from the shocked gas, we integrated the time-dependent equations for electron impact ionization and excitation through the downstream flow, using the code described by Hamilton & Sarazin (1984) and Borkowski, Sarazin, & Blondin (1994). We assumed that the H II region has the same abundances as the inner ring (Lundqvist & Fransson 1996): Z He 2.5, Z C 0.1, Z N 2, Z O 0.2, and Z Ne 0.3 (in units of solar abundance, Z A X A /X J ). We found that the fast shock model could fit the ROSAT PSPC pulse-height spectrum well, but only for Z Fe not exceeding 0.1, perhaps owing to condensation of iron into grains. (We also assumed Z Mg Z Si 0.1.) Figure 3 shows the model spectrum (including interstellar absorption). It is dominated

3 No. 1, 1997 X-RAY AND UV EMISSION FROM SNR 1987A L33 FIG. 4. UV and X-ray light curves, including interstellar extinction and absorption. From top to bottom on right-hand axis:h (dashed line); total kev X-ray flux scaled by a factor of 0.1 (solid line); N V 1240 (dashed line); O VIII Ly (dot-dashed line); N VII Ly (dotted line); O VI 1035 (dashed line); O VII He (dot-dashed line); N VI He (dotted line); He II 1640 (dashed line). by emission lines of hydrogen- and helium-like ions of N, O, and Ne. At this time, the kev X-ray energy flux is produced mainly by free-free emission (30%), N (24%), O (26%), Ne (8%), and Fe (8%). The free-free fraction increases with time. Figure 4 shows the resulting fluxes, F X (t) in the kev band and F j (t) of individual X-ray emission lines (including interstellar absorption), as functions of time. The line fluxes at first increase as EM(t) EM B (t) EM R (t) (Fig. 2) but then begin to saturate owing to the fact that C, N, O, and Ne become fully ionized and cease to emit lines farther downstream from the shocks. The condition for the flux of an emission line from a given ion, j, to saturate is n e t? [C j (T )] 1, where n e is the postshock electron density and C j (T ) is the electron impact ionization rate coefficient of ion j. For example, [C N VII (T )] cm 3 s for the 0.50 kev Ly line of hydrogenic nitrogen (where we have taken kt 2 4 kev). When a line becomes saturated, its flux will be proportional to Ṁ(t) Ṁ B (t) Ṁ R (t) (Fig. 2). Observations with AXAF will provide a critical test of this model. For example, we estimate that, with an expected 0"3 point-spread function, AXAF may be able to resolve the shock interaction into, say, four quadrants. From Figure 4, we estimate that, in 1998, the AXAF CCD imaging spectrometer will be able to detect 2700 total counts in the kev O VIII Ly line in 10 5 s. A more difficult, but extremely valuable, observation would be to resolve the line profiles with the AXAF OGS spectrometer. With a resolving power (i.e., 1200 km s 1 ) the OGS should be able to measure the line profiles of such signals, which are expected to have line widths km s 1 or more. As Luo et al. (1994) have discussed, the line profiles should depend strongly upon position and should be sensitive to the shape of the inner boundary of the H II region. 4. ULTRAVIOLET LINE EMISSION As circumstellar gas and supernova ejecta flow into the shocked zone, elements such as carbon, nitrogen, oxygen, neon, etc., are ionized to progressively higher stages by thermal electrons. At the typical (11 5 kev) electron temperatures in the shocked gas, this ionization proceeds rapidly until these elements become helium-like, then more slowly as they become hydrogenic and fully ionized. For example, the characteristic timescale for electron impact ionization of N V is given by [C N V (T )] cm 3 s. Therefore, for the model we are considering, where the shocked gas has electron density 1450 cm 3, the characteristic timescale to ionize N V is 12 weeks. The conditions in the fast shocks in SNR 1987A are similar to those in the nonradiative shocks in the remnant of SN 1006, the subject of an important study by Laming et al. (1996). They pointed out that, because lithium-like ions such as N V have much greater rate coefficients for ion impact excitation of resonance lines than for electron impact ionization, such ions will emit many ultraviolet photons before they become ionized to higher stages. Since N V ions are confined to a relatively thin layer downstream from the shock, the luminosity of N V 1240 photons produced will be proportional to the mass flux of nitrogen atoms, hence baryons, entering the shock, i.e., S 1240 t 1240 f N Ṁ t, (4) where f N is the fraction of nitrogen atoms that are less ionized than N VI. As Laming et al. (1996) show (in their Table 9), the constant 1240 increases with shock velocity and with a parameter f e characterizing the degree of initial electron-ion temperature equilibration. Here we adopt a conservative estimate, (extrapolating their value for f e 0.1 to shock velocity V s 4000 km s 1 and allowing for the elevated abundance of nitrogen). Given 1240 and f N, one can simply estimate S 1240 (t) from equation (4) and the mass fluxes shown in Figure 2. The factor f N accounts for the possibility that the nitrogen in the H II region may have been preionized to N VI or higher stages by the supernova flash (Ensman & Burrows 1992). The degree of preionization is very uncertain, however. Lundqvist & Fransson (1996) estimate that, immediately after the flash, a fraction f N (t 0) of the nitrogen at the inner surface of the ring is N V, while the rest is N VI. Even if f N (t 0) 0, it will increase to a lower limit f N (t 10 yr) as a result of radiative recombination of N VI in the H II region (assuming n II 150 amu cm 3 ). As a conservative estimate, we adopt this value. Preionization is not significant in the supernova envelope, however, which is dense enough (eq. [1]) so that the N VI can recombine very rapidly at early times (say, t 10 days), when there was no longer any radiation capable of ionizing N V. (Likewise, the currently observed X-ray flux is insufficient to preionize N V in the envelope.) Therefore, we can assume that f N 1 for the mass flux, Ṁ R (t), through the reverse shock. Figure 4 shows the light curves of N V 1240 and O VI 1035 expected for our model (including factors of 0.25 and 0.13 to account for interstellar extinction). The predicted flux of N V 1240 should be easily observable with the Hubble Space Telescope (HST); for example, we estimate a net count rate (for the entire interaction region) at t 11 yr of 24 s 1 with the Space Telescope Imaging Spectrometer (STIS) and the G140L grating. With such a count rate, it will be feasible with a long-slit spectrum to resolve the N V 1240 source both spatially and spectrally in an observation of, say, 10 4 s. Similarly, the Far Ultraviolet Spectroscopic Explorer (FUSE) should be able to observe O VI Such observations will provide a powerful diagnostic of the impact of SN 1987A with its circumstellar matter. For example, one should be able to resolve the shock interaction region

4 L34 BORKOWSKI, BLONDIN, & MCCRAY with HST much better than one can with the Australia Telescope (we expect that the N V 1240 emission-line source is coincident with the nonthermal radio source). A transverse velocity v s 5000 km s 1 corresponds to a proper motion 0"02 yr 1, which should be measurable within a few years. Moreover, as Luo et al. (1994) have illustrated, one can learn a great deal about the shock dynamics and kinetics by observing line profiles. Along the minor axis of the inner ring, the line peaks should be Doppler-shifted by the projected velocities, H3000 km s 1, of the shock interaction region. The difference of peak Doppler shifts of the blueshifted (N) and redshifted (S) images will tell the velocity and acceleration of the shock. Along the major axis, the line profiles should have complex profiles that depend on the shape of the inner boundary of the H II region. As Laming et al. (1996) have discussed, the He II 1640 emission line is excited by electrons, unlike N V 1240, which are excited by ions. Therefore, the luminosity ratio L(1640)/ L(1240) depends sensitively on the electron and ion temperatures in the gas shortly behind the shock. Detection of both lines will allow for determination of the electron/ion temperature ratio behind the reverse shock. We can also estimate that the hydrogen in the supernova envelope that is crossing the reverse shock at present should have recombined during the first few months after outburst. If so, approximately 0.3 H and4ly photons should be emitted for each hydrogen atom that crosses the reverse shock (Raymond et al. 1983). Figure 4 shows the light curve of H that we expect from this process. This line should be very broad (FWZI 2 24,000 km s 1 ), indicative of the freely expanding supernova envelope. After correction for interstellar extinction, the flux of Ly should be 14 times that of H. If these broad H and Ly emission lines are absent from the spectrum, we would infer that the hydrogen in the supernova envelope has been mostly ionized by some as yet unidentified mechanism. 5. DISCUSSION Following Chevalier & Dwarkadas (1995), we have developed a model to account for the X-ray and radio observations and to predict X-ray and UV emission-line fluxes. But many details of the model remain uncertain at this time. In particular, the H II region may have a shape very different from the fat torus that we have illustrated. Observations of the N V 1240 image and line profiles with STIS may be the best way to determine this shape. Although we developed a specific model to predict light curves of emission lines, we should emphasize that our predictions are almost model independent. That is true because the main parameters of the model the density and emission measure of the shocked gas have been set to fit the soft X-ray light curve observed by ROSAT. If the observed X-rays result from thermal emission by shocked gas, there is little latitude in adjusting the net emission measure, EM(t), or the net mass flux, Ṁ(t) 1 EM(t)/n II, through the two shocks. That is, the predicted luminosity of the N V 1240 is tied to the observed soft X-ray light curve through ratios of atomic rate coefficients. Finally, we remark on the discrepancy of the mean expansion velocity, V 2800 km s 1, of the radio source observed by Gaensler et al. (1997) and the mean shock velocity, V 4100 km s 1, in the model presented here. A slower shock would result if the H II region had a density roughly a factor of 2 greater than the value assumed here. But, if so, the emission measure of shocked gas would increase by a factor 12 3/2, resulting in an X-ray luminosity greater than that observed by ROSAT. One possible way to reconcile the discrepancy would be to assume that the H II region subtended a smaller solid angle than the fat torus assumed here. 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5 FIG. 1. Interaction of the blast wave with the H II torus at t 9 yr. The forward shock exterior to the torus is located at a radius of cm, and the reverse shock is seen along the left edge of the image. The gas with greatest density (black) is the shocked H II gas. BORKOWSKI, BLONDIN, & MCCRAY (see 476, L32) PLATE L2