arXiv:astro-ph/0203346v1 20 Mar 2002
Modelling the Hubble Space Telescope Ultraviolet and Optical
Spectrum of Spot 1 on the Circumstellar Ring of SN 1987A1
Chun S. J. Pun2,3,4, Eli Michael5, Svetozar A. Zhekov5,6, Richard McCray5, Peter M.
Garnavich7, Peter M. Challis8, Robert P. Kirshner8, E. Baron9, David Branch9, Roger A.
Chevalier10, Alexei V. Filippenko11, Claes Fransson12, Bruno Leibundgut13, Peter
Lundqvist12, Nino Panagia14, M. M. Phillips15, Brian Schmidt16, George Sonneborn3,
Nicholas B. Suntzeff17, Lifan Wang18, and J. Craig Wheeler19
1Based on 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.
2Laboratory for Astronomy and Space Physics, Code 681, NASA-GSFC, Greenbelt, MD 20771.
3National Optical Astronomical Observatories, P.O. Box 26732, Tucson, AZ 85726.
Dept. of Physics, University of Hong Kong, Pokfulam Road, Hong Kong; jc-
5JILA, University of Colorado, Boulder, CO 80309-0440.
6On leave from Space Research Institute, Sofia, Bulgaria.
7Dept. of Physics, University of Notre Dame, 225 Nieuwland Science Hall, Notre Dame, IN 46556.
8Harvard-Smithsonian Center for Astrophysics, 60 Garden St, Cambridge, MA 02138.
9Dept. of Physics and Astronomy, University of Oklahoma, 440 W. Brooks St., Norman, OK 73019.
10Dept. of Astronomy, University of Virginia, P.O. Box 3818, Charlottesville, VA 22903.
11Dept. of Astronomy, University of California, Berkeley, CA 94720-3411.
12SCFAB, Stockholm Observatory, Dept. of Astronomy, SE-10691 Stockholm, Sweden.
13European Southern Observatory, Karl-Schwarzschild-Strasse 2, D-85784 Garching, Germany.
14STScI, 3700 San Martin Drive, Baltimore, MD 21218; on assignment from the Space Science Department
15Carnegie Institution of Washington, Las Campanas Obs, Casilla 601, Chile.
16Mount Stromlo and Siding Spring Observatories, Private Bag, Weston Creek P. O., ACT 2611, Australia.
17Cerro Tololo Inter-American Observatory, Casilla 603, La Serena, Chile.
18Institute for Nuclear and Particle Astrophysics, E. O. Lawrence Berkeley National Lab, Berkeley, CA
19Dept. of Astronomy, University of Texas, Austin, TX 78712.
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We report and interpret HST/Space Telescope Imaging Spectrograph (STIS)
long-slit observations of the optical and ultraviolet (UV) (1150 − 10270˚ A)
emission-line spectra of the rapidly brightening Spot 1 on the equatorial ring
of SN 1987A between 1997 September and 1999 October (days 3869 – 4606 after
outburst). The emission is caused by radiative shocks created where the super-
nova blast wave strikes dense gas protruding inward from the equatorial ring. We
measure and tabulate line identifications, fluxes and, in some cases, line widths
and shifts. We compute flux correction factors to account for substantial inter-
stellar line absorption of several emission lines.
Nebular analysis shows that optical emission lines come from a region of cool
(Te ≈ 104K) and dense (ne ≈ 106cm−3) gas in the compressed photoionized
layer behind the radiative shock. The observed line widths indicate that only
shocks with shock velocities Vs< 250 km s−1have become radiative, while line
ratios indicate that much of the emission must have come from yet slower (Vs<∼
135 km s−1) shocks. Such slow shocks can be present only if the protrusion has
atomic density n >∼3×104cm−3, somewhat higher than that of the circumstellar
ring. We are able to fit the UV fluxes with an idealized radiative shock model
consisting of two shocks (Vs = 135 and 250 km s−1). The observed UV flux
increase with time can be explained by the increase in shock surface areas as the
blast wave overtakes more of the protrusion. The observed flux ratios of optical
to highly-ionized UV lines are greater by a factor of ∼ 2 − 3 than predictions
from the radiative shock models and we discuss the possible causes. We also
present models for the observed Hα line widths and profiles, which suggests that
a chaotic flow exists in the photoionized regions of these shocks. We discuss what
can be learned with future observations of all the spots present on the equatorial
Subject headings: supernovae: individual (SN 1987A) – supernova remnants –
Supernova 1987A (SN 1987A) in the Large Magellanic Cloud provides an unprecedented
opportunity to observe the birth and development of a supernova remnant. International
Ultraviolet Explorer (IUE) observations (Fransson et al. 1989; Sonneborn et al. 1997) found
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narrow line emission about 80 days after the explosion20, demonstrating the presence of
circumstellar gas around SN 1987A. Images taken with the Hubble Space Telescope (HST)
showed that this gas consists of an equatorial inner ring (radius ≈ 0.7 light year, ne ∼
3 × 103− 3 × 104cm−3) and two outer rings (∼ 3 times the size of the inner ring, ne <∼
2000 cm−3) tilted towards the observer at ∼ 45◦(Jakobsen et al. 1991; Wang 1991; Plait
et al. 1995; Burrows et al. 1995; Lundqvist & Fransson 1996; Maran et al. 2000; Lundqvist
& Sonneborn 2001). The circumstellar ring system was excited by the ionizing radiation from
the supernova during the shock breakout (Lundqvist & Fransson 1991; Ensman & Burrows
1992; Blinnikov et al. 2000). In the interacting winds model, the SN 1987A ring system is
part of a bipolar shell around the supernova (Blondin & Lundqvist 1993; Martin & Arnett
1995; Link et al. 2001).
The first signal of ongoing interaction between the SN 1987A debris and the circumstellar
gas was the rebirth of the supernova in X-ray (Beuermann, Brandt, & Pietsch 1994; Goren-
stein, Hughes, & Tucker 1994; Hasinger, Aschenbach, & Tr¨ umper 1996) and radio (Staveley-
Smith et al. 1992, 1993) wavelengths around day 1000. The size of the radio-emitting region
indicated that the supernova debris expanded unimpeded at velocity >∼35,000 km s−1up
to ∼ day 1000 before slowing down to ≈ 3500 km s−1by the interaction (Gaensler et al.
1997; Manchester et al. 2001). The X-ray and radio observations have been explained as the
interaction of the supernova ejecta with a rather dense (nH ≈ 100 cm−3) H II region that
separates the shocked stellar wind of the supernova progenitor from the denser gas of the
inner ring (Chevalier & Dwarkadas 1995; Borkowski et al. 1997a).
The “main event” of the birth of the Supernova Remnant 1987A (SNR 1987A) — the
interaction between the supernova debris and the circumstellar rings — has been anticipated
since the discovery of the circumstellar gas. Predictions of the time of the first contact ranged
from 2003 (Luo & McCray 1991) to 1999 ± 3 (Luo, McCray, & Slavin 1994) to 2005 ± 3
(Chevalier & Dwarkadas 1995). There have been previous model predictions of radiation
from this impact in X-rays (Suzuki et al. 1993; Masai & Nomoto 1994; Borkowski et al.
1997b) and in the UV/optical (Luo et al. 1994). The first definitive sign of impact between
the supernova blast wave and the inner ring was detected in the 1997 April HST/Space
Telescope Imaging Spectrograph (STIS) spectral images as a blueshifted (∼ −250 km s−1)
Hα feature at position angle (PA) ≈ 30◦of the ring (Sonneborn et al. 1998). Subsequent
analysis of the HST/Wide Field and Planetary Camera 2 (WFPC2) images taken in 1997
July showed that a point emission, located in projection at 88% of the distance to the ring,
was increasing in brightness over a wide range of wavelengths (Garnavich et al. 1997, 2001)
20The time of core collapse of SN 1987A, 1987 February 23.316 (UT) (JD = 2,446,849.816), was determined
by the IMB and Kamiokande II neutrino detectors (Bionta et al. 1987; Hirata et al. 1987).
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and could be traced back to as early as 1995 March (day 2932) (Lawrence et al. 2000a). The
position of the brightening spot, located just inside the edge of the inner ring, suggested that
this is the result of an interaction of the supernova blast wave with an inward protrusion
of the ring. This brightening spot, named Spot 1, has increased in flux by a factor of ∼ 10
between 1996 and 1999 (Garnavich et al. 2001). With a number of new spots (∼ 10 in 2000
November) since then detected all around the inner circumstellar ring (Garnavich et al. 2000;
Lawrence et al. 2000a), we are now witnessing the full birth of SNR 1987A.
We present here HST/STIS UV and optical spectroscopy of Spot 1 taken by the Super-
nova INtensive Study (SINS) collaboration up to day 4606 (1999 October 7). Results from
an earlier (day 3869, 1997 September 27) STIS spectrum of Spot 1 have been presented in
Michael et al. (2000). We describe the new observations and data reduction in §2 and report
the results in §3. In §2.3 we describe a detailed method to determine the intrinsic fluxes
and widths of a few UV lines, including the Si IV λλ1394, 1403 and C IV λλ1548, 1551
doublets, which are strongly affected by interstellar line absorption along the line of sight to
Our working model (Figure 1) for Spot 1, as well as the other spots, is that it is caused by
the impact of the supernova blast wave with a dense inward protrusion of the ring (Michael
et al. 2000). When the blast wave overtakes such an obstacle, slower shocks are transmitted
into it. Since a range of densities is present in the ring, we expect (§4.1) that the transmitted
shocks will have a range of velocities (Vs∼ 100 – 1000 km s−1). Not all of these shocks are
responsible for the observed UV/optical emission from Spot 1 though. While some UV and
optical line emission is produced right at the shock front, a much larger amount is produced
if the shocked gas undergoes thermal collapse, i.e. the shock becomes radiative (§4.2). The
time it takes for a shock to make this transition increases with its speed and decreases with
its pre-shock density. Once the post-shock gas collapses, a shock becomes an extremely
efficient radiator of UV and optical emission lines (§4.3). Therefore, while the range of
possible velocities present in the ring is large, we are only observing the shocks which are
slow and/or dense enough to have become radiative. Michael et al. (2000) confirmed this
general picture and found that the observed line widths and line intensity ratios indicated the
emission was formed by radiative shocks in the velocity range ∼ 100 − 250 km s−1entering
into a dense (n0>∼104cm−3) gas.
Nebular analysis of the observed emission lines of Spot 1 (§5.1) suggests that the emission
comes from a region of high density (ne> 106cm−3), This result confirms our picture that
this emission comes from radiative shocks, which can compress the pre-shock gas by a factor
>∼100. In §5.2, we model the UV line fluxes from Spot 1 between day 3869 and day 4587
with the one-dimensional steady state shock code of Cox & Raymond (1985). The observed
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UV fluxes are fitted well with models containing emission from two radiative shocks with
Vs= 135 and 250 km s−1. We propose two scenarios to interpret the observed increase of UV
flux with time of Spot 1. In the first scenario the density of the obstacle is low enough so that
the cooling times of shocks entering the obstacle are comparable to the age of the spot. The
increase in UV line fluxes are then due primarily to the aging of the shocks as they develop
thermally collapsed layers. In the second scenario, the obstacle is dense enough so that all
the shocks cool almost immediately. The increased observed fluxes are instead attributed to
an increasing surface area of shock interaction. While the second scenario fits the observed
fluxes better, we suspect that the actual light curves probably manifest a combination of both
effects. In §5.3 we discuss the observed line profiles and present simulated Hα line profiles
based on simple geometric models for the shock interaction. In §6 we discuss what we have
learned by comparing results of simplified shock models with the spectral observations, and
describe how future observations may elucidate some unsolved problems. We summarize the
main results in §7.
2.Observations and Data Reduction
The STIS observations of Spot 1 in both optical and UV wavelengths obtained by the
SINS team are listed in Table 1. SN 1987A is located in a densely populated region of the
LMC and appears to belong to a loose, young cluster region (Panagia et al. 2000). Target
acquisition was complicated by the stars present near the supernova, especially Star 3, a Be
star of V ∼ 16, at 1.′′63 away and PA = 118◦, and Star 2, a B2 III star of V = 15.0, at 2.′′91
away and PA = 318◦from the supernova (Walborn et al. 1993; Scuderi et al. 1996). We
decided to peak-up on the nearby star S2 (Walker & Suntzeff 1990) and offset the telescope
to center the aperture on Spot 1. We measured the required offset from the WFPC2 images
(Garnavich et al. 2001). Due to the uncertainties in this measurement, Spot 1 was located
at 0.′′08 off the center of the slit for all observations taken before 1999 August. We reduced
these data using the standard calibration files which assumed that the object was located at
the center of the slit. We estimate that the offset from the center of the slit will cause us to
underestimate the measured flux by <∼5% for the 0.′′2 data, and by <∼0.1% for the 0.′′5 data.
With each grating setting, we took multiple (3 − 5) observations centered at dithered
positions 0.′′5 apart along the slit. Cosmic rays (in the case of optical data) and hot pixels
(in optical and UV data) were removed simultaneously when the dithered raw images were
combined with the CALSTIS software developed by the STIS Investigation Definition Team
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at the Goddard Space Flight Center21. Previous narrow-slit STIS spectra processed by the
SINS team with this software showed that the flux calibration of far-UV (G140L) and near-
UV (G230L) data agree to <∼2% for the overlapping region, while the agreement between
near-UV (G230L) and optical (G430L) data is good to <∼5% (Baron et al. 2000; Lentz et
The location and orientation of the aperture positions are shown in Figure 2. For all
but one of the observations, the slit was oriented (within ±5◦) along the axis connecting
the center of the SN 1987A debris and Spot 1, which is located at a PA= 29◦on the inner
ring (Garnavich et al. 2001). The only exception was the 1999 October G140L observation
(Figure 2b), where the (52′′× 0.′′2) slit had a PA of 55◦and did not pass through the
center of the SN 1987A debris. In all observations, the Spot 1 spectrum overlapped with
that from the adjacent segment of the inner circumstellar ring that was included within the
STIS aperture. With an expansion velocity of 10.5 ± 0.3 km s−1(Cumming & Lundqvist
1997; Crotts & Heathcote 2000), the ring was not resolved spectrally in any of the STIS
observations reported here. In the optical data, the emission from the ring is the main
source of background to the Spot 1 spectrum and will be discussed in §2.1 and §2.2.
Garnavich et al. (2001) measured the width of Spot 1 in WFPC2 images up to 1999
April 21 (day 4440) and showed that Spot 1 was unresolved in the data and was consistent
with a point source at optical wavelengths. We compared the Full Width at Half-Maximum
(FWHM) of Spot 1 emissions in our last STIS observations at 1999 September with those of
point sources that were recorded in the data. We found that the FWHM of Spot 1 was 8±6%
and 25± 9% larger than a point source in the far-UV and optical wavelengths, respectively.
The latter result is consistent with measurements by Lawrence et al. (2000b) with the 2000
May 1 (day 4816) STIS G750M spectroscopy, which suggests that Spot 1 is now moderately
resolved in the HST data at optical wavelengths.
2.1.Low Resolution Optical Observations
We extracted the low resolution optical spectrum of Spot 1 from the STIS G430L and
G750L two-dimensional spectral images. Portions of the 1999 September G430L and G750L
data taken with the 0.′′2 slit are shown in Figures 3(c) and 3(d), respectively. The horizontal
streaks near the center in the figure are broad emission lines from the inner supernova
debris, which have a FWHM velocity vFWHM≈ 2800 km s−1(Wang et al. 1996; Chugai et
al. 1997). For these data, the lower section of the spectral image is the combined Spot 1
21CALSTIS Reference Guide, http://hires.gsfc.nasa.gov/stis/software/doc manuals.html
– 7 –
and inner ring emission-line spectrum (hereinafter referred to as the Spot 1+North-Ring, or
S1+NR, spectrum). At the kinematic resolution of the G430L and G750L grating settings
(∆V ≈ 300 − 550 km s−1), neither the emission lines from Spot 1 [vFWHM≈ 200 km s−1,
Michael et al. (2000)] or those from the ring (vFWHM≃ 10 km s−1) are resolved. The upper
section of the spectral image is the emission-line spectrum of the segment of the inner ring
subsection that is in the slit and directly opposite that of Spot 1 (hereinafter referred to as
the South-Ring, or SR, spectrum).
We measured the S1+NR and SR spectra in the G430L and G750L grating settings by
integrating the 7−9 rows of the image in which the emission-line spectra appeared. Emissions
due to the diffuse LMC background in the Balmer lines, [O II] λλ3727, 3729, and [O III]
λλ4959, 5007 are observed as images of the entire slit at those wavelengths in Figures 3(c)
and 3(d). We subtracted the contribution of this diffuse emission from the Spot 1 spectrum
by linear interpolation above and below the extracted rows. Since the subtracted LMC
background was only a small fraction of the emission from Spot 1 (<2% for [O III] λ5007),
this subtraction did not significantly increase the uncertainty of the estimated line fluxes.
We measured the flux of each emission line in the S1+NR and SR spectra by fitting
a Gaussian to the line profile, allowing the net flux, width, and central wavelength to vary
independently for each line. We estimated the background level from a linear fit over a 45-
pixel region of the spectrum centered on the line in question but excluding any emission-line
features. We fitted the line flux by minimizing the total χ2in which the coefficients defining
the line and the background were free parameters. Gaussian profiles gave satisfactory fits to
all the line profiles. The reduced-χ2, or χ2
of freedom), of our line fits were within the range 0.8 − 2.8, compared to the value of 1.0
for a statistically good fit. We computed the flux of each line and its error from the best-fit
parameters and their associated uncertainties. We adjusted the statistical errors of all the
line fluxes so that χ2
in wavelength, such as [O III] λ5007 + He I λ5016, and [Ar III] λ7136 + [Fe II] λλ7155, 7172,
we fitted the emission features with multiple Gaussians with the additional constraint that
their wavelength separations were the known differences of laboratory wavelengths.
r, (the total χ2divided by the number of degrees
r= 1.0 for all fits. In cases where two or more emission lines overlapped
Only a subset (∼ 1/3) of the emission lines observed in the S1+NR spectrum also
appeared in the SR spectrum. Therefore, we attributed entirely to Spot 1 the measured
fluxes of emission lines that were seen in the S1+NR spectrum but not in the SR one.
To subtract the NR spectrum from the S1+NR spectrum, we assumed that the fluxes of
emission lines in the NR spectrum are equal to those in SR spectrum scaled by factors that
are independent of time. This assumption is reasonable because the rate of flux decrease
around the ring has been shown to be relatively constant around the ring (Plait et al. 1995;
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Lundqvist & Sonneborn 2001). We estimated the scale factors by examining archival WFPC2
emission-line images in Hα, [O III] λ5007, and [N II] λ6583 obtained in 1994 February and
1994 September, before Spot 1 appeared. We measured flux ratios f(NR)/f(SR) of 1.2,
1.1, and 1.3, respectively from the Hα, [O III] λ5007, and [N II] λ6583 images. We used
the Hα scale factor for the Balmer lines. While we used the [O III] λ5007 factor for the
[O III] λλ4363, 4959, 5007 lines, and likewise for the [N II] lines, we recognized the increase
in systematic errors in the measured Spot 1 fluxes of [O III] λ4363 and [N II] λ5755 because
these lines are more temperature sensitive than the other lines. For all remaining emission
lines, such as [S II] and [Ne III], we assumed a scale factor of 1.2 ± 0.1 in order to subtract
the NR spectrum.
In addition to the 0.′′2 data sets described above, we obtained one observation of Spot 1
with the 0.′′5 slit and the G750L grating setting. As before, we extracted the combined
S1+NR spectrum by integrating the 6 rows of the image where the emission lines appeared.
For the emission lines produced predominantly by Spot 1, such as He I λ6678 and [Ar III]
λ7135, we measured the S1 fluxes by fitting the line emissions with single Gaussian profiles.
For the lines where emission from both S1+NR were apparent, such as [N II] λ5755 and
[O I] λ6300, we estimated the contribution of the NR flux to the S1+NR flux from a linear
interpolation of the NR emission-line flux adjacent to Spot 1 in the slit. After subtracting
this background, we fitted the remaining S1 emission lines with Gaussian profiles. The
uncertainties in the background ring flux estimated in this procedure resulted in larger
systematic errors in the estimated Spot 1 line fluxes for this observation.
2.2. Medium Resolution Optical Observations
A medium resolution optical spectrum was taken on 1999 August 30 (UT, 4368.0 days
since explosion) with a 0.′′1 slit and the G750M (6581) grating setting (6295−6867˚ A). With a
spectral resolution of ∆V ≃ 50 km s−1, emission lines from Spot 1, with vFWHM≈ 200 km s−1
(Michael et al. 2000), were resolved in the data, while the emission from the unshocked inner
ring, with vFWHM≃ 10 km s−1, remained unresolved. Three observations of 7800 seconds
each were taken at three parallel slit positions, pointed so that the middle slit position
was centered on Spot 1 and the two other slit positions were immediately adjacent [cf.
Figure 2(d)]. Therefore the observation covered a segment of the ring of length 0.′′3.
With the crowded stellar field near SN 1987A, we did not execute the acquisition-peakup
exposure for these 0.′′1 slit observations as suggested by the HST/STIS operation manual
(Leitherer et al. 2000). Instead all three adjacent slit positions were placed on Spot 1 by
blind offsets. Therefore we cannot apply the standard pipeline data reduction procedures to
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process these data. To determine the fluxes of emission lines from Spot 1, we first removed
the wavelength-dependent aperture throughput correction function in the pipeline data for
each 0.′′1 slit spectrum. We then summed the fluxes from the three 0.′′1 slit positions, and
multiplied the total flux by a new aperture correction function for an equivalent 0.′′3 slit,
calculated by interpolating the pipeline corrections for the 0.′′1, 0.′′2, and 0.′′5 slits. For Hα,
the aperture correction led to a ≃ 20% decrease of flux over the sum of the fluxes measured
in the three slit positions. We calculated the corresponding 1σ errors by combining the
individual errors in quadrature.
A section of the spectral image from the middle slit position is shown in Figure 3(e).
Again, the central horizontal streak is emission from the SN 1987A debris. Emissions from
the inner ring at the two positions where the ring intersected with the slit aperture were
observed in [O I] λλ6300, 6364, Hα, [N II] λλ6548, 6583, and [S II] λλ6717, 6731. In
the lower section of the spectral image, emission from the stationary inner ring overlapped
with the broadened emission from Spot 1. Again, the major source of contamination the
Spot 1 spectrum is the emission from the inner circumstellar ring within the 0.′′1 slit. As
described by Michael et al. (2000), we fit all the Spot 1 emission features with Gaussian
functions. Emission from the stationary ring dominated the spectral profile near zero velocity
(∼ ±50 km s−1) and was excluded from the fit. The majority of the line profiles could be
fitted well with Gaussian profiles, such as the fit to the [N II] λ6583 line emission shown in
Figure 4. The sole exception was the Hα line profile, where the signal was strong enough to
show noticeable departures from a Gaussian profile, as we shall discuss further in §5.3.
2.3. Low Resolution UV Observations
We obtained low resolution UV spectra of Spot 1 with the G140L and G230L grating
settings. Michael et al. (2000) have already presented results from the first G140L far-UV
observations in 1997 September 27 taken with the 0.′′5 slit. We detected emission lines from
Spot 1 in N V λλ1239, 1243, Si IV λλ1394, 1403, O IV] λ1400, N IV] λλ1483, 1487, C IV
λλ1548, 1551, and He II λ1640. We detected the same set of emission lines in 1999 February
27, observing again with the 0.′′5 slit. In 1999 October 7, observing with the 0.′′2 slit, we also
detected the C II λ1335 multiplet, [Ne IV] λ1602, and O III] λλ1661, 1666. Figure 3(a) shows
a section of the spectral image from this observation. Radiation from Spot 1 is evident in the
lower portion of the image whereas only faint line emission from the inner circumstellar ring
can be seen in the upper half of the displayed image. The broad (∼ ±15,000 km s−1) Lyα
radiation comes from the reverse shock from the interaction between the supernova debris
and the H II region located inside of the equatorial ring (Sonneborn et al. 1998; Michael
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et al. 1998a,b). Fluxes of UV line emission from the inner ring are much less than those
from Spot 1 and make a negligible contribution to the measured fluxes. This is also the case
for the near-UV emission lines measured in the 1999 September 17 G230L 0.′′2 observation,
shown in Figure 3(b).
We measured the far-UV and near-UV spectra of Spot 1 from the G140L and G230L
data, respectively, by procedures similar to those we described in §2.1. We fitted emission
lines with Gaussian profiles and linear backgrounds, except for the broad Lyα emission
underlying the N V λλ1239, 1243 doublet which we fitted with a quadratic function. We
fitted the two components of close doublets such as N V λλ1239, 1243, C IV λλ1548, 1551,
O III] λλ1661, 1666, N II] λλ2139, 2143, and Mg II λλ2796, 2803 with Gaussians constrained
to have fixed doublet separations, identical widths, and line ratios dictated by their oscillator
strengths. At the resolution of the G140L grating setting, the Si IV λ1403 emission of the
Si IV λλ1394, 1403 doublet is blended with the O IV] λ1400 multiplets. We fit the combined
Si IV and O IV] feature with multiple Gaussians, requiring the Si IV doublet to meet the
same constraints as the other close doublets.
The observed fluxes of a few UV lines, such as Si IV λλ1394, 1403 and C IV λλ1548,
1551, were reduced by interstellar line absorption. We describe our procedure for correcting
for this absorption in §3.3 below.
3.1. Optical Emission Line Fluxes
Table 2 lists the measured Spot 1 optical emission line fluxes, including previously
published results from the 1998 March 7 (day 4030) data by Michael et al. (2000) and 1σ
upper flux limits for the [O II] λλ3726, 3729, and [N I] λλ5198, 5200 doublets. The 1σ errors
tabulated are only statistical errors. Systematic effects, such as fringing for the G750L data
towards the near-IR region (λ > 8500˚ A, Leitherer et al. 2000), might contribute additional
uncertainties to the measured fluxes.
The tabulated fluxes have been dereddened with E(B −V ) of 0.16 (Fitzpatrick & Wal-
born 1990) and the extinction law of Cardelli et al. (1989) with an assumed RV of 3.1. In the
optical band the differences between the LMC extinction law and the Galactic law are neg-
ligible at low color excess (Fitzpatrick 1999). The interstellar extinction correction applied
is listed in the last column of Table 2.
Spot 1 was observed in many neutral and lowly ionized species in the optical wavelengths.
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We did not detect any coronal lines such as [Fe X] λ6375. Most of the line fluxes increased
with time at a rate comparable to that measured from WFPC2 photometry (Garnavich et
al. 2001). At the low spectral resolution of both the G430L and G750L observing modes
(R = λ/∆λ = 530 − 1040), definitive line identification remained a problem, especially for
the [Fe II] emission lines. The Fe line identifications in Table 2 are based upon the modeling
of the Spot 1 Fe lines in Pun et al. (2002, in preparation). Moreover, several lines were
blended. Table 2 lists possible contributing species and, in the case of [Fe II] lines, different
multiplets. In contrast, line blending is not a problem in the medium resolution G750M
(R ≈ 6000) data.
Fluxes in [O I] λλ6300, 6363 and He I λ6678 were measured with the low resolution
G750L and medium resolution G750M gratings in 1999 September within 17 days of each
other. The two measurements agreed within uncertainties for the [O I] doublet. The two
results differed by ∼ 50% for the low S/N He I λ6678 data, but were also within the noise
level. This difference is probably a good indication of the detection limits of such faint lines.
3.2. Optical Emission Line Widths
For the medium resolution G750M observations, apart from the line fluxes, we were also
able to measure the peak emission velocities (V0) and the widths (VFWHM) of the emission
lines from the profile fits. Table 3 lists the results. The peak emission velocity measurements
have been adjusted for the SN 1987A heliocentric velocity of +286 km s−1(Wampler et al.
1989). The 1998 March 0.′′2 G750M results have also been adjusted for the off-center position
of Spot 1 within the aperture (cf. §2). The main uncertainties in the measurements of both V0
and VFWHMare due to the contributions by emission from the stationary circumstellar ring,
which dominated the emission by Spot 1 near zero velocity for many species (cf. Figure 4).
The errors due to this contribution are generally smaller for the 1999 August 0.′′1 observations
than the 1998 March 0.′′2 ones.
For all emission lines, the line centroids from Spot 1 were blueshifted, with peak velocity
V0 lying within the range −40 to −10 km s−1. This result is consistent with the overall
physical picture in which Spot 1 is located on the near side of the equatorial ring (Sonneborn
et al. 1998) and the shock entering Spot 1 is moving towards the observer. We found no
evidence that the peak velocity V0of Spot 1 varied with time.
Most emission lines had widths (FWHM) within the range ∼ 150− 180 km s−1, except
[N II] λ6583 and Hα, which had a slightly greater width, VFWHM>∼200 km s−1. We detected
no measurable change with time of the line widths except for Hα and [N II] λ6548. For Hα,
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the emission profile from the second observation in 1999 August could no longer be fit well by
a single Gaussian (to be discussed further in §5.3). The width of the [N II] λ6548 emission,
measured in 1999 August was 1.5 times greater than that measured in 1998. However, we
are inclined to attribute this increase to systematic error, since we detected no such increase
in the other [N II] component at 6583˚ A, where the line fluxes were ∼ 3 times stronger.
3.3.UV Emission Lines
3.3.1. Interstellar Line Absorption
Near the time of outburst, interstellar absorption lines of C II λ1335 multiplet, Si IV
λλ1394, 1403, C IV λλ1548, 1551, and Mg II λλ2796, 2803 were observed in the UV con-
tinuum of SN 1987A with IUE operating in the high resolution (∆V ≈ 30 km s−1) echelle
mode (Blades et al. 1988; Welty et al. 1999). For each line, the dominant absorption compo-
nent was centered at +281 km s−1, near the SN 1987A heliocentric velocity of +286 km s−1
(Wampler et al. 1989), and had a FWHM of ∼ 80 km s−1. Therefore, narrow emission lines
(vFWHM≃ 10 km s−1) from the inner circumstellar ring from these species are completely
blocked by the interstellar absorption (Fransson et al. 1989), as demonstrated by the absence
of Si IV and C IV emission lines from the ring in Figure 3(a).
Emission lines from Spot 1, with VFWHM≈ 200 km s−1, are not totally blocked by these
interstellar absorption lines, as correctly predicted by Luo et al. (1994). However, the line
profiles are altered and the observed fluxes are reduced substantially. Therefore, we must
correct the measured C II, Si IV, C IV, and Mg II emission line fluxes to account for this
absorption. The appropriate correction factors depend on the profile shapes of both the
Spot 1 emission and the intervening absorption. To make this correction, we assumed that
the UV emission lines from Spot 1 had Gaussian profiles with the same parameters as the
optical lines as measured in the medium resolution optical data (§2.2). We then estimated
the amount of flux reduction in the UV emission lines by multiplying the assumed Gaussian
profiles by the absorption profiles seen in the IUE data. The corresponding flux correction
factors for the various emission lines are shown in Figure 5(a). The correction factor is
greatest for the Mg II λλ2796, 2803 doublet, where the interstellar absorption is almost
completely saturated between −50 and +300 km s−1.
The flux correction factor due to interstellar absorption is sensitive to the assumed
widths (VFWHM) and peak (V0) velocities of assumed line profiles from Spot 1. Figure 5(b)
illustrates this dependence for the important case of C IV λλ1548, 1551. We see that the
flux correction factor increases moderately with decreasing VFWHMfor VFWHM> 200 km s−1,
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but becomes very sensitive to both VFWHMand V0for VFWHM< 150 km s−1.
As we shall discuss below in §5.3, the peak velocities and widths of the emission lines
from Spot 1 depend on the detailed geometry and hydrodynamics of the shocks entering the
spot, which are unknown. It is not obvious that the optical and UV emission lines should
have the same peak velocities and widths. However, in the plane-parallel shock model that
we describe in §4, we found that the peak velocities of emission from Spot 1 were almost
identical in both the UV and optical wavelengths. Therefore, we used the measured peak
Spot 1 velocities from the medium resolution optical observations, V0= −30 ± 15 km s−1
(cf. Table 3), to estimate the flux correction factors to account for interstellar absorption of
the UV emission lines.
3.3.2. Widths of the UV Emission Lines
The far-UV emission lines from Spot 1 are poorly resolved at the spectral resolution of
the G140L grating (∆V ≈ 250 km s−1at 1500˚ A). For these lines, we attempted to establish
the relation between the actual and observed widths through Monte Carlo simulations. For
the simulations, we assumed that the intrinsic Spot 1 emission lines had (i) a Gaussian shape
with FWHM as a variable parameter, (ii) emission peaked at −30 km s−1, and (iii) flux ratio
of the doublets dictated by their oscillator strengths.
The dotted curve in Figure 6(a) shows a model C IV λλ1548, 1551 profile, assuming
VFWHM = 150 km s−1and line flux ratio I(1548)/I(1551) = 2:1. The solid curve shows
a subsection of the high resolution IUE spectrum with the absorption line profile for the
C IV doublet. For each input FWHM, we multiplied the model input emission line by the
measured IUE absorption spectrum. Figure 6(b) shows a typical result. We convolved this
profile with the line-spread function (LSF) of the detector for the slit aperture of the data
set. The LSF for each emission line was interpolated from the measured LSFs at 1200˚ A,
1500˚ A, and 1600˚ A. The LSF typically has a narrow peak (∼ 1.5 pixel) and a broad wing
that extends to ∼ 10 pixels on either side of the peak (Leitherer et al. 2000). Figure 6(c)
shows a typical example of the LSF for the G140L grating at 1550˚ A with the 0.′′2 aperture.
For each emission line, we normalized the resulting convolved profile to the photon
counts for each observation. The thick solid curve in Figure 6(d) shows the normalized
model C IV line profile for the 1999 October observation. We constructed a simulated
observed profile (the thin solid curve in Figure 6d) by applying random noise to the model
profile and sampling it at the spectral resolution of the G140L grating. Finally, we fitted
the simulated profile with Gaussians in the same way as the real data. The dotted curve in
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Figure 6(d) shows such a fit.
For each input model FWHM velocity, we ran 10000 simulations and generated an array
of the corresponding observed line widths. Figure 7 shows the median and the 68.3% (1σ)
upper and lower limits of the array plotted against the model input widths of the various
observed emission lines. We use these results to convert the observed line widths from the
STIS data, shown as horizontal lines in Figure 7, to the intrinsic widths of the lines emitted
by Spot 1.
We did not attempt to model the near-UV G230L data in this way because the spectral
resolution for this grating setting (∆V ≈ 450 km s−1at 2500˚ A) was too low for such a
procedure to yield useful results.
For all emission lines except the N V doublet, we found a 1σ upper limit of FWHM <∼
300 km s−1, consistent with the measurements from the medium resolution optical data.
The lower limit to the assumed FWHM is important because the flux correction factors for
all UV emission lines are sensitive to the input line widths for FWHM < 150 km s−1. We
took this lower limit to be the same as the measured lower limit for the optical lines, that
is, FWHM > 100 km s−1.
3.3.3.UV Line Fluxes and Line Widths
Table 4 lists the fluxes of UV lines from Spot 1 inferred from the Gaussian line fits (cor-
rected for extinction) and the intrinsic line widths derived from the Monte Carlo simulations.
The tabulated fluxes for the C II, Si IV, C IV, and Mg II doublet have been corrected for
interstellar line absorption assuming an intrinsic Spot 1 line width VFWHM= 150+150
with peak emission at V0= −30 ± 15 km s−1. The flux correction factors applied for these
lines are also listed in Table 4. We also corrected the fluxes for interstellar extinction by
assuming E(B −V )LMC= 0.06, E(B −V )Galactic= 0.10, and RV = 3.1. At UV wavelengths,
the choice of extinction function is important because the correction is substantial and is
known to vary from place to place (cf. Pun et al. 1995). We used the 30 Doradus extinction
function of Fitzpatrick (1986) for the LMC component, and the Seaton (1979) function for
the Galactic component.
The far-UV G140L data had slightly higher kinematic resolution (∆V ≈ 250 km s−1)
than the low resolution optical G430L and G750L data. However, the individual components
of C II λ1335 and O IV]λ1400 multiplets remained unresolved in the data. There were more
uncertainties with line identifications in the lower resolution (∆V ≈ 300 − 650 km s−1)
near-UV G230L data, such as the unidentified emission features near 2737˚ A and 2746˚ A.
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On the other hand, we identified emission features near 2324˚ A and 2334˚ A as the C II
λ2325 multiplet and Si II λ2335, respectively. We ruled out the alternative identification
of [O III] λλ2322, 2332 because the observed line ratios, I(2324˚ A)/I(2334˚ A) ≈ 5 and
I(2324˚ A)/I(4363˚ A) ≈ 7, were much different from the theoretical [O III] line ratios of
∼ 280 and 0.12, respectively. These [O III] line ratios are determined only by the atomic
transition probabilities and are independent of the excitation model.
The fluxes of all far-UV lines from Spot 1 increased with time during the three obser-
vations taken from 1997 September (day 3869) to 1999 October (day 4606). The rate of
increase differed for features from different ions, ranging from I(3869 d)/I(4596 d) = 5.3
for N V λ1240 to 1.9 for C IV λ1550. We will discuss the time dependence of the UV
emission-line fluxes in §5.2.
4.1. Impact Hydrodynamics
In our working model for Spot 1, we assume that the UV and optical emission lines
observed are caused by radiative shocks that develop where the supernova blast wave strikes
dense gas protruding inward from the circumstellar ring. As we shall show, the spectrum
and profiles of the emission lines from such shocks are sensitive to the density distribution
and geometry of this protrusion, which are probably quite complex and cannot be resolved
even with the HST. Our approach here therefore is to illustrate the salient physics of the
spectrum formation with a few idealized “toy models,” which we believe will guide us toward
a better understanding of the properties of a more realistic model.
Following the previous work of Luo et al. (1994) and Borkowski et al. (1997b), we show
in Figure 8 hydrodynamic simulations for two models of a fast shock overtaking a dense gas
cloud. In each model, we assume that a fast plane-parallel blast wave traveling through a
uniform medium of relatively low density overtakes a cloud of substantially greater uniform
density. The cloud boundary is approximated as a density discontinuity. In each simulation,
the blast wave drives a transmitted shock into the cloud, while a reflected shock travels
backwards towards the interior of the remnant. As the blast wave overtakes the obstacle,
the surface area of the transmitted shock increases. The transmitted shock propagates with
a range of velocities depending on shape and density of the obstacle.
The simulations are calculated using the hydrodynamic code VH-1 (Strickland & Blondin
1995), which is based on the piecewise parabolic method of Colella & Woodward (1984). Ra-
diative cooling is included in the code using an operator splitting technique. We have used