The Astrophysical Journal, 592:1137–1150, 2003 August 1 # 2003. The American Astronomical Society. All rights reserved. Printed in U.S.A.
HUBBLE SPACE TELESCOPE SPECTROSCOPY OF THE UNEXPECTED 2001 JULY OUTBURST OF THE DWARF NOVA WZ SAGITTAE1 Edward M. Sion Department of Astronomy and Astrophysics, Villanova University, 800 Lancaster Avenue, Villanova, PA 19085;
[email protected]
Boris T. Ga¨nsicke Department of Physics and Astronomy, University of Southampton, Southampton SO17 1BJ, UK;
[email protected]
Knox S. Long Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218;
[email protected]
Paula Szkody Department of Astronomy, University of Washington, Seattle, WA 98195;
[email protected]
Fu-Hua Cheng Center for Astrophysics, University of Science and Technology of China, Hefei, Anhui 230026, People’s Republic of China
Steve B. Howell Institute of Geophysics and Planetary Physics, University of California, Riverside, CA 92521;
[email protected]
Patrick Godon2 Department of Astronomy and Astrophysics, Villanova University, 800 Lancaster Avenue, Villanova, PA 19085;
[email protected]
William F. Welsh Department of Astronomy, San Diego State University, San Diego, CA 92182-1221;
[email protected]
Sumner Starrfield Department of Physics and Astronomy, Arizona State University, Tempe, AZ 85287; sumner.starrfi
[email protected]
Christian Knigge Department of Physics and Astronomy, University of Southampton, Southampton SO17 1BJ, UK;
[email protected]
and Warren M. Sparks XNH, MS F664, Los Alamos National Laboratory, Los Alamos, NM 87545;
[email protected] Received 2002 November 25; accepted 2003 April 17
ABSTRACT We present Hubble Space Telescope (HST) Space Telescope Imaging Spectrograph E140M spectra of the dwarf nova WZ Sge, following the early superoutburst of 2001 July. Our four far-ultraviolet (FUV) spectra, obtained over a time span of 4 months, monitor changes in the hot component of the system during the ˚ . They reveal Stark-broadened Ly decline phase. The spectra cover the wavelength interval 1150–1708 A and He ii (1640) absorption and absorption lines due to metals (Si, C, N, Al) from a range of ionization states. Single-temperature white dwarf models provide reasonable qualitative agreement with the HST spectra. We find that the white dwarf appears to dominate the spectra from October through December. However, it is not clear that the absorption lines of metals form in the white dwarf photosphere. Therefore, the derived abundances and rotational velocity must be viewed with caution. Only a modest improvement in the fits to the data results when an accretion belt component is included. If the FUV spectra arise from the white dwarf alone, then we measure a cooling in response to the outburst of at least 11,000 K (29,000–18,000 K). The absence of broad underlying absorption features due to metals at this stage suggests slow rotation (200 km s1). It is possible that the white dwarf envelope has expanded due to the heating by the outburst or that the relatively narrow absorption features we observe are forming in an inflated disk atmosphere or curtain associated with the outburst. Subject headings: novae, cataclysmic variables — stars: individual (WZ Sagittae) — white dwarfs orbital periods, longest outburst recurrence times, lowest mass Roche lobe–filling secondaries, lowest accretion rates, and coolest white dwarf primaries of any class of dwarf novae (Howell et al. 1999, 2002; see Kato et al. 2001 for a recent review). It is also the brightest dwarf nova during outburst and arguably the closest cataclysmic variable, with a distance of only 43 8 pc (J. Thorstensen 2001, private communication). The inclination of the binary is high enough (78 ) that the secondary star eclipses the disk rim but not the white dwarf. Steeghs et al. (2001) found the radial velocity semiamplitude, K2 , of the secondary star to
1. INTRODUCTION
WZ Sge is the widely known, extensively studied prototype of a group of H-rich cataclysmic variables that have extreme properties: the largest outburst amplitudes, shortest 1 On the basis of observations with the NASA/ESA Hubble Space Telescope, obtained at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555. 2 Visiting at the Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218.
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be in the range 493–595 km s1, giving an upper limit to the mass of the secondary, M2 < 0:11 M , while they found an upper limit to the radial velocity semiamplitude of the primary, K1 < 37 km s1, implying a mass range 0:7 M < Mwd < 0:9 M from the mass function. A 28 s periodicity was detected in both the optical (Patterson 2002) and the far-ultraviolet (FUV; Welsh et al. 1997). If due to rotation, this oscillation period would correspond to V sin i ¼ 1200 km s1 for a 1 M white dwarf. The outburst is widely held to arise from an extended period of enhanced mass accretion onto the white dwarf, during which the accretion disk is expected to be luminous. This disk radiates the FUV spectrum during the outburst phase, but as the accretion rate declines, the disk luminosity also falls, and the white dwarf emerges as the dominant source of FUV radiation. Further flux declines are attributed to the cooling of the white dwarf photosphere, which was heated by the outburst. Following the 1978 December outburst, its relatively clocklike outburst regularity of every 33 yr was unexpectedly disrupted by a 10 yr outburst on 2001 July 23, first reported by T. Ohshima (Ishioka et al. 2001). What followed was the most thoroughly observed dwarf nova outburst in history (Patterson et al. 2002; Kuulkers et al. 2002; Knigge et al. 2000). Shortly following the discovery of the superoutburst, two director’s discretionary proposals for HST observation were approved, one (GO-9287) to provide dense coverage of
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the outburst phase and the other (GO-9304) to cover, with a preliminary analysis, the emergence of the accretion-heated white dwarf and monitor its evolution (cooling and other properties) in response to the superoutburst. The results of the latter program, awarded four orbits over a 4 month time span, are reported in this paper. 2. THE HST OBSERVATIONS
The observations took place on 2001 September 11 (03:03:28 UT; 50 days postoutburst), October 10 (12:57:13 UT; 79 days postoutburst), November 10 (11:45:32 UT; 110 days postoutburst), and December 11 (00:34:59 UT; 141 days postoutburst) with HST Space Telescope Imaging Spectrograph (STIS) using the FUV MAMA detector configuration in time-tag mode with the E140M grating and 0>2 0>2 aperture, giving a wavelength coverage of 1140– ˚ centered on 1425 A ˚ . We binned the echelle spectra in 1735 A ˚ steps and analyzed the average STIS spectra from 0.1 A each individual HST visit. The data reduction was carried out with the standard STScI pipeline reduction system. The spectral line measurements were carried out with the software SPLOT in IRAF by fitting Gaussians. The results of the line identifications and line strength measurements are presented in Tables 1–4, where, in each, column (1) lists the centroid wavelength, column (2) lists the central depth line flux in ergs cm2 s1, column (3) lists the equivalent width
TABLE 1 Line Identifications and Absorption Line Measurements, 2001 September Spectrum Line Center ˚) (A (1)
Flux ˚ 1) (ergs cm2 s1 A (2)
EW ˚) (A (3)
FWHM ˚) (A (4)
Line Identification (5)
1175.04 .................... 1191.27 .................... 1193.71 .................... 1197.12 .................... 1200.38 .................... 1206.21 .................... 1228.29 .................... 1240.52 .................... 1250.55 .................... 1253.82 .................... 1260.33 .................... 1264.71 .................... 1300.50 .................... 1309.26 .................... 1317.21 .................... 1329.24 .................... 1335.09 .................... 1369.86 .................... 1393.31 .................... 1401.70 .................... 1434.60 .................... 1493.06 .................... 1526.25 .................... 1532.78 .................... 1548.58 .................... 1640.49 .................... 1658.12 .................... 1671.31 .................... 1701.84 .................... 1719.60 .................... 1724.25 ....................
3.8 1012 1.7 1013 5.3 1013 1.4 1013 3.6 1013 5.0 1013 2.4 1013 1.7 1012 1.6 1013 1.3 1013 9.3 1013 1.3 1012 7.0 1012 2.7 1013 1.5 1013 1.1 1013 1.7 1012 2.2 1013 1.7 1012 1.6 1012 4.7 1013 2.9 1013 3.5 1013 3.2 1013 2.0 1012 1.3 1012 2.6 1013 5.7 1013 2.3 1013 2.1 1013 2.3 1013
6.0 0.5 1.6 0.5 1.3 2.1 0.7 4.4 0.3 0.2 1.6 2.3 11.6 0.6 0.2 0.2 3.5 0.4 4.0 3.9 1.1 0.7 1.0 0.9 5.2 4.6 1.0 2.1 0.8 0.9 1.0
6.7 1.0 2.0 1.1 2.1 2.7 3.5 6.5 1.4 1.0 2.3 3.6 13.8 1.3 1.5 1.6 3.7 3.0 5.9 5.9 5.6 5.4 2.1 1.9 6.6 14.3 4.7 3.5 4.0 3.3 3.4
C iii 1175.26, 1175.59, 1175.71, 1175.99 Si ii 1190.41 Si ii 1193.29, 1194.50 Si ii 1197.39 N i 1200.22, 1200.71 Si iii 1206.50 Si ii 1228.43, 1228.61; Si ii 1228.74 N v 1240 S ii 1250.50 S ii 1253.79 Si ii 1260.42 Si ii 1264.73, 1265.00 Si iii 1298.89, 1298.96, 1301.15; Si iii 1303.32 Si ii 1309.27 Ni ii 1317.22 C i 1329.09, 1329.10, 1329.2 C ii 1334.53, 1335.66, 1335.71 Ni ii 1370.13 Si iv 1393.76; Ni ii 1393.32 Si iv 1402.77 ? N i 1492.03, 1492.82, 1484.68 Si ii 1526.70 Si ii 1533.43 C iv He ii 1640.33, 1640.35, 1640.39, 1640.47 C i 1657.91, 1658.12 Al ii 1670.85 or Fe ii 1673.46 Fe ii 1701.94, 1702.04 Al ii 1719.49 + Fe ii Al ii 1725.11
TABLE 2 Line Identifications and Absorption Line Measurements, 2001 October Spectrum Line Center ˚) (A (1)
Flux ˚ 1) (ergs cm2 s1 A (2)
EW ˚) (A (3)
FWHM ˚) (A (4)
Line Identification (5)
1152.03 .................... 1156.58 .................... 1164.48 .................... 1168.17 .................... 1175.58 .................... 1190.04 .................... 1193.54 .................... 1197.01 .................... 1200.11 .................... 1206.04 .................... 1228.23 .................... 1238.64 .................... 1242.79 .................... 1250.45 .................... 1253.67 .................... 1260.30 .................... 1264.81 .................... 1272.14 .................... 1302.62 .................... 1309.30 .................... 1317.09 .................... 1329.27 .................... 1335.34 .................... 1370.28 .................... 1393.73 .................... 1402.77 .................... 1435.12 .................... 1466.74 .................... 1492.26 .................... 1494.16 .................... 1526.33 .................... 1533.05 .................... 1549.30 .................... 1658.10 .................... 1671.05 .................... 1673.67 .................... 1701.60 .................... 1719.78 .................... 1724.11 ....................
4.9 1013 3.3 1013 1.5 1013 2.6 1013 1.0 1012 1.5 1013 2.4 1013 7.9 1014 1.7 1013 1.8 1013 1.4 1013 2.6 1013 2.8e 1013 1.4 1013 1.5 1013 8.2 1013 9.4 1013 1.2 1013 4.1 1012 3.3 1013 1.1 1013 8.4 1014 1.3 1012 1.5 1013 8.7 1013 7.1 1013 3.8 1013 2.4 1013 1.0 1013 8.9 1014 3.6 1013 4.5 1013 1.5 1012 2.3 1013 3.3 1013 2.5 1013 2.0 1013 1.9 1013 1.8 1013
1.298 0.8 0.4 0.7 3.2 1.0 2.1 0.9 1.9 2.3 1.9 1.6 1.4 0.6 0.5 2.5 3.1 0.3 10.6 1.1 0.3 0.2 3.8 0.4 2.5 2.2 1.2 0.8 0.4 0.3 1.5 1.8 5.0 1.3 1.7 1.3 1.1 1.1 1.1
2.419 2.7 1.7 3.3 3.3 1.3 2.3 1.1 2.6 3.1 4.3 1.9 1.7 1.2 1.1 2.7 3.4 1.8 12.8 1.4 1.4 1.8 3.7 2.3 2.8 2.5 9.2 7.5 1.2 1.3 2.0 2.2 5.7 4.5 1.9 5.0 3.6 2.6 3.1
O i 1152.15 P ii 1156.96 N ii 1164.28, 1164.58 N i 1167.45, 1168.54 C iii 1175.26, 1175.59, 1175.71, 1175.99 Si ii 1190.41 Si ii 1193.29, 1194.50 Si ii 1197.39 N i 1200.22, 1200.71 Si iii 1206.50 Si ii 1228.74, 1228.43, 1228.61 N v 1240 N v 1240 S ii 1250.50 S ii 1253.79 Si ii 1260.42 Si ii 1264.74 Fe ii ? Si iii 1298.89, 1298.96, 1301.15; Si iii 1303.32, O i 1302.17 Si ii 1309.27 Ni ii 1317.22 C i 1329.09, 1329.10,1329.2 C ii 1334.53, 1335.66, 1335.71 Ni ii 1370.13 Si iv 1393.76; Ni ii 1393.32 Si iv 1402.77 ? Ni ii 1467.26 N i 1492.03, 1492.82 N i 1494.68 Si ii 1526.70 Si ii 1533.43 C iv C i 1657.91, 1658.12 Al ii 1670.85 Fe ii 1673.46 Fe ii 1701.94, 1702.04 Al ii 1719.49 + Fe ii Al ii 1725.05, 1725.07, 1725.11
TABLE 3 Line Identifications and Absorption Line Measurements, 2001 November Spectrum Line Center ˚) (A (1)
Flux ˚ 1) (ergs cm2 s1 A (2)
EW ˚) (A (3)
FWHM ˚) (A (4)
1175.45 .................... 1189.66 .................... 1193.27 .................... 1196.68 .................... 1199.51 .................... 1206.14 .................... 1238.44 .................... 1242.48 .................... 1250.34 .................... 1253.52 .................... 1260.19 .................... 1264.58 .................... 1296.53 .................... 1298.74 ....................
5.0 1013 1.1 1013 2.1 1013 7.0 1014 1.5 1013 1.8 1013 1.6 1013 5.9 1014 6.0 1014 6.3 1014 4.4 1013 7.0 1013 6.0 1014 1.3 1013
2.6 1.1 2.7 1.3 2.2 3.1 1.7 0.4 0.3 0.3 2.2 3.0 0.3 0.7
2.6 1.7 2.7 1.8 3.3 4.0 2.1 1.2 1.2 1.0 2.4 3.6 1.1 1.1
Line Identification (5) C iii 1175.26, 1175.59m, 1175.71, 1175.99 Si ii 1190.41 Si ii 1193.29, 1194.50 Si ii 1197.39 N i 1200.22, 1200.71 Si iii 1206.50 N v 1240 N v 1240 Si ii 1250.50 S ii 1253.79 Si ii 1260.42 Si ii 1264.74 Si iii 1296.73 Ti iii 1298.70
TABLE 3—Continued Line Center ˚) (A (1)
Flux ˚ 1) (ergs cm2 s1 A (2)
EW ˚) (A (3)
FWHM ˚) (A (4)
Line Identification (5)
1301.88 .................... 1304.25 .................... 1309.10 .................... 1317.01 .................... 1329.27 .................... 1335.06 .................... 1393.40 .................... 1402.48 .................... 1433.62 .................... 1438.71 .................... 1465.15 .................... 1492.99 .................... 1494.16 .................... 1526.01 .................... 1532.78 .................... 1548.79 .................... 1657.88 .................... 1670.91 .................... 1719.98 .................... 1723.79 ....................
1.0 1013 8.5 1014 2.0 1013 1.3 1013 8.4 1014 1.0 1012 6.1 1013 5.1 1013 2.8 1013 2.2 1013 4.7 1013 1.5 1013 8.9 1014 2.8 1013 4.0 1013 1.2 1012 2.6 1013 3.8 1013 9.8 1014 1.5 1013
0.7 0.7 0.9 0.2 0.2 4.2 2.3 2.0 1.1 0.9 2.0 0.7 0.3 1.4 1.9 5.0 1.6 1.8 0.7 1.2
1.5 1.2 1.3 1.4 1.7 4.1 2.5 2.3 7.9 4.1 12.9 5.3 1.3 1.9 2.5 5.4 5.2 2.0 2.1 3.4
Si iii 1298.89, 1298.96, 1301.15; Si iii 1303.32; O i 1302.17 Si ii 1304.37 Si ii 1309.27 Ni ii 1317.22 C i 1329.09, 1329.10, 1329.2 C ii 1334.53, 1335.66, 1335.71 Si iv 1393.76; Ni ii 1393.32 Si iv 1402.77 S i 1433.30; Ca ii 1433.75 ? Fe ii 1465.04 N i 1492.03, 1492.82 N i 1494.68 Si ii 1526.70 Si ii 1533.43 C iv C i 1657.91, 1658.12 Al ii 1670.85 Al ii 1719.49 + Fe ii Al ii 1725.05, 1725.07, 1725.11
TABLE 4 Line Identifications and Absorption Line Measurements, 2001 December Spectrum Line Center ˚) (A (1)
Flux ˚ 1) (ergs cm2 s1 A (2)
EW ˚) (A (3)
FWHM ˚) (A (4)
Line Identification (5)
1175.73 .................... 1190.51 .................... 1193.76 .................... 1197.22 .................... 1200.10 .................... 1206.29 .................... 1238.78 .................... 1242.05 .................... 1250.58 .................... 1253.82 .................... 1260.44 .................... 1264.85 .................... 1294.64 .................... 1296.90 .................... 1299.05 .................... 1302.15 .................... 1304.60 .................... 1306.06 .................... 1309.46 .................... 1317.15 .................... 1335.43 .................... 1393.86 .................... 1402.89 .................... 1418.86 .................... 1492.91 .................... 1526.46 .................... 1533.21 .................... 1549.28 .................... 1658.08 .................... 1671.29 .................... 1673.67 ....................
4.8 1013 6.5 1014 7.9 1014 3.3 1014 4.7 1014 7.1 1014 1.2 1013 1.2 1013 4.8 1014 5.2 1014 2.6 1013 3.8 1013 5.0 1014 4.4 1014 1.0 1013 5.5 1014 6.0 1014 1.6 1014 1.6 1013 7.3 1014 8.0 1013 4.7 1013 4.0 1013 9.9 1014 1.6 1013 1.8 1013 3.1 1013 8.7 1013 1.4 1013 1.7 1013 2.5 1013
3.1 1.1 1.8 0.9 1.2 1.8 1.7 1.4 0.4 0.4 1.9 2.6 0.3 0.3 0.7 0.5 0.6 0.9 0.9 0.3 3.8 2.3 1.9 0.5 0.9 1.2 1.8 4.4 1.2 1.4 1.3
3.0 1.4 2.1 1.4 1.7 2.1 2.1 1.8 1.5 0.8 2.1 3.2 0.7 1.0 1.2 1.1 1.1 1.3 1.3 1.7 3.8 2.4 2.2 3.1 5.7 1.5 2.5 5.4 3.8 1.5 5.0
C iii 1175.26, 1175.59, 1175.71, 1175.99 Si ii 1190.41 Si ii 1193.29, 1194.50 Si ii 1197.39 N i 1200.22, 1200.71 Si iii 1206.50 N v 1240 N v 1240 S ii 1250.50 S ii 1253.79 Si ii 1260.42 Si ii 1264.74 Si iii 1294.55; Ti iii 1294.67; Ti iii 1294.72 Si iii 1296.73 Si iii 1298.89, 1298.95 O i 1302.17 Si ii 1304.37 O i 1306.03 Si ii 1309.27 Ni ii 1317.22 C ii 1334.53, 1335.66, 1335.71 Si iv 1393.76; Ni ii 1393.32 Si iv 1402.77 ? N i 1492.03, 1492.82 Si ii 1526.70 Si ii 1533.43 C iv C i 1657.91, 1658.12 Al ii 1670.85 Fe ii 1673.46
OUTBURST OF WZ SGE ˚ ), column (4) lists the FWHM (in A ˚ ), and column (EW in A (5) gives the line identification(s). The analysis of the FUV light curves obtained in the time-tag mode will be reported elsewhere. The September observation took place during a rebrightening, as shown in Figure 1 of Long et al. (2003), where the placement of the Far Ultraviolet Spectrographic Explorer (FUSE) and HST STIS observations is compared with the optical light curve of the outburst and decline. This rebrightening was one of several closely spaced so-called echo outbursts. These echo outbursts, smaller in amplitude and shorter in duration than the main outburst, are widely thought to be smaller accretion events. The exact physical mechanism responsible for the echo outbursts is not yet understood. The spectroscopic results are discussed in the next section. 3. THE SPECTRAL FEATURES
The four spectra corresponding to the September, October, November and December HST observations are displayed as flux versus wavelength plots in Figures 1–4, together with the line identifications labeled in the figures
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with vertical tick marks. First, there is a steady drop in the ˚ ) from 4:2 1013 continuum flux level (measured at 1450 A ˚ 1 in September to 1:9 1013 ergs cm2 s1 ergs cm2 s1 A ˚ 1 in December, suggesting that the white dwarf is cooling A and/or shrinking after being bloated by the outburst heating. Although it would be useful to compare this continuum flux level decline, there is an unfortunate gap in the IUE coverage of the postoutburst interval following the 1978 December 1 superoutburst. Therefore, there are no IUE spectra for direct comparison with our spectra at a comparable epoch in time. The earliest postoutburst spectrum following the 1978 outburst was taken a full 7 months later, whereas our 2001 December spectrum was taken 5 months ˚, 7 postoutburst. The flux level of SWP5761 at 1450 A ˚ 1 months postoutburst onset, is 1:6 1013 ergs cm2 s1 A compared with the flux level of our 2001 December HST ˚ 1, 5 months spectrum of 1:9 1013 ergs cm2 s1 A postoutburst onset. The most prominent line feature in all four spectra is the very broad Ly absorption, which we attribute to the high-gravity white dwarf photosphere. One does not see the H2 quasi-molecular absorption feature (centered around
˚ 1) vs. wavelength (A ˚ ) for the 2001 September 11 observation of the early superoutburst of WZ Sge. Vertical tick Fig. 1.—Observed flux, F (ergs cm2 s1 A marks: Strongest line features.
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Fig. 2.—Same as Fig. 1, except for the 2001 October 10 observation
˚ ) in any of the four spectra, although this feature is 1400 A prominent during quiescence when the temperature of the white dwarf is below about 20,000 K (Sion et al. 1995c). The other features in Figures 1–4 are predominantly in absorption, except for apparent broad emission wings at C iv (1548, 1550) flanking deep absorption. There is, however, a weak, double-peaked Ly feature in emission that could be associated, at least in part, with the system. A close examination reveals that the feature appears doublepeaked, especially in the November and December data, and that the profile resembles that of C iv in that the blue emission peak is higher than the red peak for October– December. In the September data, it is more difficult to see, but it looks plausible that the red peak is higher, just as it is in C iv. A rough measurement of the peak-to-peak separa˚ . At Ly, 1 A ˚ corretion of the peaks is approximately 6 A sponds to about 245 km s1, so that separation corresponds to roughly 1500 km s1. This leads to a value of the disk velocity, Vdisk sin i ’ 750 km s1. For comparison, Skidmore et al. (2000) measure Vdisk sin i ¼ 630 km s1 in quiescence from the Balmer H. Therefore, these values for Ly suggest a disk-formed line. This is extremely important because it bears directly upon the basic question of the reestablishment of the disk following the outburst. On the other
hand, a chromosphere (temperature inversion) on a rapidly rotating white dwarf might produce lines with these widths. The other line features cover a broad range of ionization. Of immediate interest is the mix of ions, the excitation/ ionization states of the transitions, and the differences one sees between these spectra and the spectra of WZ Sge obtained during quiescence. First, the presence of N v in absorption in the four spectra is extremely important since it was never seen in spectra obtained during quiescence. For example, there is no evidence of the N v (nor C iii) absorption in the lower quality, lower resolution IUE spectra from early 1979 July. It is noteworthy here that N v is seen in absorption in the dwarf nova U Gem during quiescence but does not share the same velocity as the gravitationally redshifted white dwarf photosphere (Long et al. 1994). This feature is thought to arise in an extended hot region of gas near the white dwarf in U Gem. Moreover, the N v feature in WZ Sge is too sharp to be consistent with formation in a velocity-broadened disk region at the high inclination of WZ Sge unless it arises from absorption in the outer disk of continuum photons from the inner disk and/or white dwarf. ˚ that has an interesting behavThere is a feature at 1608 A ior: the absorption becomes stronger with time. While the feature is likely identified as Fe ii, other identifications are
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Fig. 3.—Same as Fig. 1, except for the 2001 November 11 observation
possible but less likely (e.g., N v, C iii). We note the presence of artifacts of STIS in the form of pseudoabsorptions at ˚. approximately 1652, 1671, 1690, and 1710 A There are noteworthy temporal variations seen in the line strengths. A number of line features became increasingly narrow with time, and emission wings present in September are not seen in December. The most striking example of this change in profile shape is the Si iv doublet. It is quite possible that the Si iv feature in the September spectrum may be associated with an accretion belt on the white dwarf spun up by the high accretion that took place during the outburst. In that spectrum, the red wing of the C iv doublet compo˚ exhibits P Cygni emission. However, by the nent at 1550 A time of the October spectrum and continuing through the December spectrum, it is the blue wing of C iv 1548 that appears in emission. While double-peaked profile structure is typically associated with formation in the accretion disk during dwarf nova quiescence, the feature may possibly signal a major change in the outflow characteristics and therefore the emergence of the underlying accreting white dwarf accompanied, however, by signs of disk accretion. It is also ˚ ) is seen in broad, shallow interesting that He ii (1640 A absorption in the September spectrum but only marginally present in the October, November, and December spectra.
This feature would not be expected to form in the photosphere of a 20,000–30,000K white dwarf. The weakening of He ii with time is presumably indicative of declining temperature in its formation region. Finally, numerous absorption lines of neutral, singly and doubly ionized sulfur, silicon, and carbon are present. Neutral nitrogen lines also appear. Given the distance to WZ Sge of only 43 pc, it is unlikely these features are, in any significant measure, interstellar in origin because they appear to be too broad for interstellar features. Some of the low-ionization features, e.g., Si ii 1260, 1265, 1527, and 1533 ˚ , have similar shapes and widths to N v, suggesting the A possibility that these low-ionization features form in the same nonphotospheric region. Other low-ionization lines, ˚ , do not resemble expected photospheric feae.g., C ii 1335 A tures. The predominance of absorption lines from neutral carbon one sees in deep quiescence spectra is not found in the four spectra. Nevertheless, the white dwarf spectrum has emerged recognizably by the time of the September observation (see Fig. 1), on the basis of our comparative examination of the archived 2001 August spectra of Knigge et al. (2002). There are also unidentified features between 1430 and ˚ present, to some degree, in all four spectra, which we 1440 A
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Fig. 4.—Same as Fig. 1, except for the 2001 December 10 observation
tentatively identify as a blend of Si i and C i. Such features were seen in the quiescence spectra of WZ Sge (Sion et al. 1995; Cheng et al. 1997). It is noteworthy that the disk emission seen in quiescence (Sion et al. 1995c; Cheng et al. 1997) when the disk is optically thin is not seen at all in the four spectra (September, October, November, and December) following the 2001 July outburst. However, it is possible that the accretion disk is optically thick in the lines but optically thin in the continuum once it is reestablished following the outburst. As we discuss below, we cannot conclusively rule out an optically thick, steady state accretion disk.
dwarf model spectra are normalized to 1 R at a distance of 1 kpc, the distance of paffiffiffiffi source is computed from d ¼ 1000 ðpcÞ ðRwd =R Þ= S , where S is the scale factor S ¼ 4ðRwd =R Þ2 ðd=kpcÞ2 . In all cases, we first found the best-fitting model and then computed the stellar radius from the scale factor by using the distance of 43 pc. Before the model fitting, any emission lines, including those in the core of Ly , were masked out in each individual STIS spectrum. In Table 5, we provide the results for accretion disk–only fits to the four spectra. This exercise utilized model accretion disks for the full range of white dwarf masses (0.35–1.2 M ), accretion rates (108.5 to 1010:5 M yr1), and disk inclination angles (18 –81 ) presented in the grid of Wade &
4. SYNTHETIC SPECTRAL FITTING
The model atmosphere (TLUSTY; Hubeny 1988), and spectrum synthesis (SYNSPEC; Hubeny, Lanz, & Jeffrey 1994; Hubeny & Lanz 1995) codes and details of our 2 minimization fitting procedures are discussed in detail in Sion et al. (1995a, 1995b) and Huang et al. (1996a, 1995b) and will not be repeated here. To estimate physical parameters, we took the white dwarf photospheric temperature Teff , log g, Si and C abundances, and rotational velocity (Vrot ) as free parameters and computed the reduced 2 . Since our white
TABLE 5 Accretion Disk Only
Observation
_ M ðM yr1)
i (deg)
2
Scale Factor
Mwd ðM Þ
Sep.................... Oct.................... Nov................... Dec ...................
109.0 109.0 109.5 109.5
18 18 18 18
13.11 16.92 12.13 7.87
1.93 101 3.17 101 1.19 100 9.54 101
0.55 0.35 0.35 0.35
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TABLE 6 White Dwarf Only
Observation
T (1000 K)
Si (solar)
C (solar)
V sin i (km s1)
2
Scale Factor
log g (cm s2)
Sep............................ Oct............................ Nov........................... Dec ...........................
29.0 +0.5/0.5 22.0 +0.8/1.0 19.0 +0.2/0.1 18.0 +0.1/0.1
5.0 +0.3/0.3 1.0 +0.1/0.1 0.5 +0.1/0.1 0.5 +0.1/0.1
2.0 +3.0/0.4 0.5 +0.1/0.1 0.1 +0.2/0.1 0.1 +0.1/0.2
600 +100/100 200 +50/50 200 +50/50 200 +20/50
8.43 10.57 7.63 5.10
2.71 102 5.00 102 7.14 102 7.33 102
8.5 8.5 8.5 8.5
Note.—Error bars are 3 .
Hubeny (1998), except for the models at inclinations of 8 , which have numerical problems. These fits were carried out without regard to any observed constraints on the white dwarf mass, disk inclination, or accretion rate. For each of the four HST observations, the best-fitting disk model was determined, and the results are presented in Table 5. The best-fitting disk model in each case was for either a white dwarf mass of 0.4 or 0.35 M , an accretion rate between _ ¼ 108:5 and 109:5 M yr1, and disk inclination angles M of i ¼ 18 . It is clear that the best-fitting disk models, selected without regard to observed parameters, are implausible because the observed absorption lines are far too sharp and therefore must arise from another source. The white dwarf mass used in the best-fitting disk models is also implausibly low, especially in view of the results reported by Steeghs et al. (2001). In addition, the best-fitting disk inclination angle is far too low to be applicable to WZ Sge, which, in quiescence, suffers hot spot eclipses and has an inclination close to 80 . This conclusion is consistent with the results of Howell et al. (1999, 2002), who demonstrated that no inner disk or boundary layer is present during the superoutbursts of tremendous outburst amplitude dwarf novae (TOADS). However, it is still possible that a model with a plausible inclination and white dwarf mass could give a reasonable fit to the spectra even if, in a numerical sense, the reduced 2 value is not optimal. Since the Wade & Hubeny disk models are normalized to 100 pc, then the distance to a source in parsecs is simply the 100 pc divided by the square root of the scale factor. Using the distance of 43 pc as a constraint on the fits, there are three models that provide roughly the correct flux. These models, all at a disk inclination of 75=5, are models P, T, and Y in the nomenclature of Table 2 in Wade & Hubeny (1998). They have white dwarf masses of 0.8, 1.0, and 1.2 M , respectively, and accretion rates, respectively, _ ¼ 109:5 , 1010.0, and 1010:5 M yr1. Thus, by conof M straining the fits to satisfy the distance to WZ Sge and by adopting the most reasonable estimates of the white dwarf mass and binary inclination, viable values of the accretion rate are derived that do not depend on the details of continuum slope fitting. In an additional exercise, we kept the white dwarf mass fixed at 1.2 M , the inclination at 81 , and fitted disk models to the October spectrum, letting the accretion rate vary. The best-fitting disk model (lowest reduced 2 ) corresponded to _ ¼ 1010:5 M yr1. This is considerably lower than the M accretion rate implied by the best-fitting disk model _ ¼ 109:0 M yr1) in Table 5, where fits using all (M combinations of the disk parameters were carried out. Our next experiment utilized single-temperature white dwarf–only fits to the HST data. In these fits, the white
dwarf temperature was varied from 16,000 to 30,000 K in steps of 1000 K, the surface gravity was chosen to be log g ¼ 8:5, rotational velocity V sin i varied from 200 to 800 km s1 in steps of 100 km s1, and chemical abundances of Si and C were chosen to be either 0.01, 0.1, 1, 2, 3, 5, or 10 times the solar values. The resulting best-fit parameters and their 3 errors for each of the four spectra are recapitulated in Table 6. The best-fitting single-temperature, log g ¼ 8:5, white dwarf models to the September, October, November, and December observations are displayed in Figures 5, 6, 7, and 8, respectively. If log g ¼ 8:0 is used in the fitting, the temperatures in the above table become 28,000 K for the September spectrum, 20,000 K for the October spectrum, 18,000 K for the November spectrum, and 17,000 K for the December spectrum. These temperatures are 800–1600 K cooler than those derived at higher gravity. As a check on our temperature results from scaling the models to the observations, we have computed an extensive grid of models in fine-temperature steps of 100 K and interpolated within this grid for log g ¼ 8:5. The resulting temperatures for the four spectra are 28,800 K (September), 21,600 K (October), 19,200 K (November), and 18,300 K (December). The single-temperature white dwarf fits indicate that following the outburst, the accretion-heated white dwarf cooled by at least 11,000 K (29,000–18,000 K). In Figure 9, we display a cooling curve of temperature versus time with the days counted starting on the first day of the outburst. In the figure, we have denoted the STIS data points with circles and the FUSE data points (from Long et al. 2003) with stars. The rotational velocity, V sin i, from the singletemperature white dwarf models, on the basis of the fitting of the relatively narrow metal features and Ly profile, is only 200 km s1. This velocity, if it corresponds to the true underlying white dwarf, is far lower than the value of 1200 km s1 reported by Cheng et al. (1997) during quiescence. We suspect, however, that these narrow features may not be forming in the white dwarf photosphere (see below). We calculated the white dwarf radius from the scale factors listed in Table 6 by fixing the distance at 43 pc. The resulting white dwarf radii corresponding to the bestfitting single-temperature white dwarf models to the four spectra are 4:92 108 (September), 6:68 108 (October), 8:00 108 (November), and 8:10 108 cm (December). The December value of the white dwarf radius implies a white dwarf mass smaller than 0.6 M . Such a low mass is ruled out by the radial velocity study of Steeghs et al. (2001). We also tried various combinations of white dwarf models and disk models as composite fits. All but one such fit
Fig. 5.—September HST STIS spectrum of WZ Sge compared with the best-fitting single-temperature white dwarf model atmosphere: Teff ¼ 29; 000 K, log g ¼ 8:5, Si = 5 solar, C = 2 solar, and V sin i ¼ 600 km s1.
Fig. 6.—October HST STIS spectrum of WZ Sge compared with the best-fitting single-temperature white dwarf model atmosphere: Teff ¼ 22; 900 K, log g ¼ 8:5, Si = 1 solar, C = 0.5 solar, and V sin i ¼ 200 km s1.
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Fig. 7.—November HST STIS spectrum of WZ Sge (flux vs. wavelength) compared with the best-fitting single-temperature white dwarf model atmosphere: Teff ¼ 19; 000 K, log g ¼ 8:5, Si = 0.5 solar, C = 0.1 solar, and V sin i ¼ 200 km s1.
Fig. 8.—December HST STIS spectrum of WZ Sge compared with the best-fitting single-temperature white dwarf model atmosphere: Teff ¼ 18; 000 K, log g ¼ 8:5, Si = 0.5 solar, C = 0.1 solar, and V sin i ¼ 200 km s1.
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Fig. 9.—Cooling curve of temperature vs. time for the WZ Sge white dwarf with temperatures from HST STIS and FUSE measurements. Time in days is counted starting from the first day of the outburst. Circles: STIS data points. Stars: FUSE data points.
failed with large 2 values. The September spectrum could be fitted with a white dwarf having Teff ¼ 28; 000 K, log g ¼ 8:5 with Si elevated to 5 times solar, C with 2 times solar plus an accretion disk with Mwd ¼ 0:80 M , with _ ¼ 1010:5 M yr1, and i ¼ 41 . This best-fit disk plus M white dwarf combination is displayed in Figure 10. The dot-
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ted line represents the white dwarf model, the dashed line is the accretion disk contribution, and the solid line is the combination of the two components. Note that the accretion disk contributes less than 1% of the FUV flux in this combined fit. However, this model lacks plausibility because we know the inclination of the disk in WZ Sge is close to 80 . In an additional experiment, two-temperature, differentially rotating composite fits were attempted, consisting of a rapidly spinning equatorial accretion belt at higher temperature and a cooler, slowly rotating photosphere. For example, for the November spectrum, the best-fitting two-temperature combination corresponded to an accretion belt with V sin i ¼ 3400 km s1, Tbelt ¼ 28; 000 K, log g ¼ 6, and solar abundances, combined with an 18,000 K photosphere with log g ¼ 8:5, V sin i ¼ 200 km s1, Si = 0.5 solar, and C = 0.1 solar. Overall, however, these composite fits produced only a very modest improvement in the reduced 2 value. For the November observation, the best-fitting white dwarf plus accretion belt corresponds to 88% of the flux emitted by the white dwarf and 12% emitted by the accretion belt. Finally, we have attempted to assess the effect of a layer of cool gas in which iron peak absorbers would alter the emergent stellar spectrum. In this model, an absorbing layer of gas is located between the star, and the observer like a curtain (which is why this model is referred to as the iron curtain). We adopted curtain parameters similar to those found in the FUSE study of WZ Sge by Long et al. (2003). We adopted Tcur ¼ 10; 000 K, electron density in number per cubic centimeter Ne ¼ 1013 , vturb ¼ 200 km s1, and
Fig. 10.—September HST STIS spectrum of WZ Sge compared with the best-fitting single-temperature white dwarf model atmosphere plus accretion disk model. The white dwarf has Teff ¼ 28; 000 K, log g ¼ 8:5, Si = 5 solar, C = 2 solar (dotted line) and the accretion disk has Mwd ¼ 0:80 M , M ¼ 1010:5 M yr1, disk inclination angle i ¼ 41 (dashed line). Solid line: Sum of the disk plus the white dwarf spectrum.
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hydrogen column density log NH ¼ 20:24. However, with the use of the curtain option in SYNSPEC (adopting Tcur and electron density, computing the mass density , constructing a curtain opacity table, and using the SYNSPEC curtain algorithm CIRCUS with input hvturb i and NH for an adopted fractional coverage of the stellar source), the resulting spectrum differs very little from the synthetic spectrum without a curtain and makes little difference in the quality of the fit, at least in the STIS range. If a column is chosen with log NH > 22, the theoretical spectrum begins to deviate significantly from the observed spectrum. However, we find that for a two-component absorbing layer consisting of a cool iron curtain with T ¼ 10; 000 K and a column density of NH ¼ 1022 plus a warmer iron curtain with T ¼ 30; 000 K and a column density of NH ¼ 1022 , a better fit to the details of absorption lines results, but not especially to the overall shape of the continuum. In particular, the fits to the lines C iii 1175, Si ii 1190, 1193, 1197, Si iii 1206, Si ii 1260, 1265, Si ii 1309, C ii 1334, Si iv 1394, 1403, and Si ii 1527, 1533 are improved. We also tried a slowly rotating (200 km s1) white dwarf model and a fast-rotating (1200 km s1) white dwarf model, both with a two-component high column density absorbing layer (with a turbulent velocity of 200 km s1), and did not notice any appreciable difference. In some parts of the spectrum, the fast-rotating star model seems to fit the observation slightly better, but in other regions, it is the slowly rotating models with curtains that appear to fit better.
5. DISCUSSION
We have analyzed the STIS data with fits of synthetic spectra of optically thick accretion disks, high-gravity photospheres, combined disks and photospheres and twotemperature, differentially rotating, composite fits combining photospheres with rapidly spinning accretion belts. An optically thick accretion disk added to a white dwarf photosphere does not satisfactorily match the observations for the range of accretion disk models that were available to us. Single-temperature white dwarf models do provide reasonable agreement with the HST spectra, but the September spectrum is problematic to this interpretation. It appears that there is an additional radiating component. For example, the Ly line does not approach zero flux as one expects for a 20,000 K white dwarf. Perhaps there is filling in of the absorption by some light from the disk. Another point is that, unless only a fraction of the white dwarf is exposed in the September data, then an increasing white dwarf radius with time is indicated as we progress to the December data. The system was observed during an optical rebrightening phase in which an additional source of radiation would be expected in the September data. It seems clear that our interpretation of overall changes in the white dwarf (e.g., cooling, radius changes) hinges critically on what radiating components are contributing to the September spectrum. Among the possible sources are a circumbinary disk, ejected shell, or corona/wind structure on the scale of several white dwarf radii, although in the latter two sources, one would expect signs of line emission that are not seen. Since Chandra observations revealed a forest of emission lines just after outburst (Kuulkers et al. 2002), this may be evidence
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of a large, optically thin component extending out to a few white dwarf radii. The problems with applying a white dwarf photospheric interpretation alone to the September spectrum may also extend to the other three spectra as well. We cannot as yet be certain that the narrow absorption lines of metals in all four spectra really form in the white dwarf photosphere. Therefore, the abundances and rotational velocity we have derived must be regarded as preliminary. The main points in favor of a white dwarf interpretation of the FUV spectra are (1) that it is difficult to produce the broad, observed Ly absorption feature with anything other than a high-gravity, hydrogen-rich photosphere and (2) that the model fits give reasonable white dwarf radii for the known distance (43 pc) of WZ Sge. It is possible that the underlying white dwarf remains obscured, perhaps due to a raised disk rim during the outburst that may gradually become less vertically extended as the outburst progresses, thus leading to the exposure of an increasing fraction of the white dwarf photosphere. The obscuration of the white dwarf by the occulting effect of a disk rim in a high-inclination cataclysmic variable has been demonstrated recently in an HST study of the nova-like variable DW UMa by Knigge et al. (2000). The presence of material high above the disk plane of WZ Sge is also suggested by the X-ray data of Kuulkers et al. (2002). If the FUV spectra arise from the white dwarf alone, then we measure a cooling in response to the outburst from 30,000 to 19,000 K. This is a lower limit because we did not observe the peak temperature the white dwarf reached during the outburst. Even if the first observation has to be excluded until a better understanding of the disk contribution to that spectrum is obtained, the white dwarf cooled by almost 5000 K between October and December. The absence of the broad absorption features (as seen in deep quiescence) suggests that the white dwarf has expanded due to the heating by the outburst or that the relatively narrow absorption features we observe are forming in an inflated disk atmosphere or curtain of remnant material associated with the outburst. In a companion paper (Long et al. 2003), we performed a similar analysis of the FUSE spectra obtained in WZ Sge in a similar period. On the whole, the two analyses agree, although the temperatures derived from the FUSE data are somewhat higher than those obtained here. It is not entirely clear whether the higher FUSE temperatures represent a nonuniform temperature distribution on the surface of the white dwarf or some effect associated with the narrow lines in the spectrum. We are hopeful that analysis of spectra obtained far from the outburst with FUSE and HST will resolve these inconsistencies as well as resolve the overall question of the rotation rate of the white dwarf. This work is supported by NASA through grants GO9304 from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555. Support was also provided, in part, by NSF grant 99-01195 and NASA ADP grant NAG5-8388 (E. M. S.). B. T. G. acknowledges support from a PPARC Advanced Fellowship.
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