The X-Ray-resolved Supernova Remnant S8 in the Dwarf Irregular Galaxy IC 1613

Nearby galaxies are natural subjects of deep Chandra observations because of Chandra's high angular resolution (90% encircled energy within 1''). Chandra observations have dramatically increased the number of known discrete X-ray sources in nearby galaxies (e.g., Long et al. 2010 for M33; Holt et al. 2003 for NGC 6946; among many others). The high accuracy of the positions of detected X-ray sources permits searches for counterparts at other wavelengths.

Supernova remnants (SNRs) represent one end stage for stellar evolution of massive stars undergoing core collapse or white dwarfs pushed past the Chandrasekhar mass limit. Through a broad display of explosion physics, SNRs potentially provide insight into the explosion mechanism(s). Vink (2012) provides a recent review of X-ray emission from Milky Way (MW) SNRs.

Extragalactic SNRs are increasingly playing a role in broadening the range of parameters beyond the MW examples (e.g., Lacey et al. 1997; Matonick & Fesen 1997; Pannuti et al. 2007; Long et al. 2010). SNRs are difficult to study in the MW because their distances are often uncertain and the intervening column of material can obscure portions of their spectra. The study of extragalactic SNRs is particularly important for those host galaxies that are not spirals, or for which the metallicity is very different, or the star formation rate differs, among additional parameters. Furthermore, distances are a much smaller problem because the distance to the host galaxy is often known to a high degree. Leonidaki (2017) provides a recent review of this developing area.

In this paper, we describe a Chandra observation that resolves an SNR in the nearby irregular Local Group galaxy IC 1613. E. M. Schlegel et al. (2019, in preparation) describe the overall results from the IC 1613 observation; this paper was extracted separately to bring attention to the SNR. The SNR, first noted by Sandage (1971) as an H ii region, was listed in his paper as S8—we adopt that label throughout this paper. Additional details about S8 are described in Section 3.

IC 1613 is classified as a DDO Irr V (van den Bergh 2000) lying within the Local Group. Estimates of its distance land around ∼730 kpc (Dolphin et al. 2001; Karachentsev et al. 2004). The most recent distance determination is 724 ± 17 kpc (Hatt et al. 2017). Here we adopt 725 kpc. IC 1613 is known to be a low-metallicity galaxy ([Fe/H] = −1.3; Kunth & Östlin 2000). The known Galactic column density, N_H, in the direction of IC 1613 is ∼2 × 10²⁰ cm⁻² based on the observed E_B−V ∼ 0.025 (Schlafly & Finkbeiner 2011) and the low-metallicity relation of Fitzpatrick (1985). Given that the galaxy is a dwarf, its massive star formation history is of interest. This is particularly important for IC 1613 given the very massive OB association in the northeast quadrant of the galaxy and the relative lack of star clusters elsewhere (Wyder et al. 2000).

Prior observations of S8 are briefly reviewed in Section 3.1. In the X-ray band, Lozinskaya et al. (1998) used the ROSAT High Resolution Imager (HRI) to observe IC 1613 twice for a duration of ∼20 ks each. Four sources were observed: a background galaxy cluster, the SNR S8, an m_v 11.4 foreground star, and a probable X-ray binary. Sources 1, 3, and 4 of that list are described in E. M. Schlegel et al. (2019, in preparation); here we describe the second source.

We used the back-illuminated S3 CCD of Chandra's Advanced CCD Imaging Spectrometer (ACIS; Garmire et al. 2003) to observe IC 1613 on 2005 September 4 (obs ID: 5905) for an effective exposure of 49327 s. The Chandra observation covers roughly the central one-half of the galaxy. The star formation/bubble complex in the northeast quadrant of IC 1613 was the target of the observation. The SNR lies south of the center of the complex and within the region of the CCD with the sharpest focus.

The Chandra Interactive Analysis of Observations (CIAO) software (version 4.9) and the associated calibration files (version 4.7) were used in the analysis (Fruscione et al. 2006). To test for the presence of soft background flares, we accumulated a source-free background area offset south from the galaxy center and away from the galaxy cluster. We extracted a light curve using 50 s bins; no flares were detected.

The SNR events were extracted in two ways: to obtain a spectrum and to obtain its color–color information. For the spectrum, we extracted the counts in the 0.3–8 keV band from a region surrounding the SNR with a radius of 45 and centered at R.A. 01:05:02.4, decl. +02:08:42.1. A source-free region from the same CCD and $0\buildrel{\,\prime}\over{.} 5$ south of the SNR was extracted for the background. Response matrices and effective area files were separately generated for the SNR and the background.

Additional details of the observation and the data reduction are presented in E. M. Schlegel et al. (2019, in preparation). Those details are mostly related to the other sources detected in IC 1613 and are not useful for discussions of S8.

3.1. Overview of Past Observations

3.2. Chandra Image

Figure 1 shows a tiling of the Chandra image, the Chandra image with contours, an Hα image, and a merged Swift Ultraviolet Optical Telescope (UVOT) NUV image (filter UVW1) of the immediate surroundings of S8. From the UV image, it is clear that S8 exists in a region of UV-bright, hence massive, stars. S8 is resolved in the Chandra image: the diameter is ∼55, which corresponds to ≈19 pc at the distance to IC 1613. Rosado et al. (2001) infer a diameter of 24 pc for S8 from Hα emission—that the Hα diameter is larger could be attributed to cooling gas and snow-plowed matter. Regardless, either diameter places S8 on the small side of the SNRs in M33 (Long et al. 2010) or the LMC (Ou et al. 2018), for example.

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Visually, the X-ray image of S8 resembles a nearly complete circle with slightly enhanced emission to the south. There is a hint of a single high, approximately centered pixel in both images, possibly a point source (position: 01:05:02.44, +02:08:41.9). Alternatively, that location could be a spot of significant interaction in the ejecta. A stronger conclusion is not possible at present because of the distance to S8 and Chandra's resolution. That could change with the detection of a compact object, for example. The Chandra image then suggests that S8 could be a plerion. We return to that identification in Section 5.

Lozinskaya et al. (1998) described S8 as resembling the LMC SNR N 49. In comparison with SNR catalogs (e.g., Williams et al. 1999), however, we think a more appropriate comparison lies with N 63A (e.g., Warren et al. 2003): in N 49, there is a strong interaction in one direction (southeast) with a significant decrease in emission in the opposite direction. For S8, there is enhanced emission to the south, but the drop in the opposite direction is less steep—a description that mimics N 63A. Rosado et al. (2001) also make the comparison between S8 and N 63A.

Figure 2 shows a hardness image of S8 where the hardness ratio (HR) is defined as $\mathrm{HR}=\tfrac{H-S}{H+S}$ and where the bands are defined as S = 0.4–0.75 keV and H = 0.75–2 keV. These energy ranges were chosen to maximize the available spectral information based on the observed spectrum (Section 4): they separate low-Z lines (N, O) from higher-Z lines (Ne, Si, Fe L) typically found in SNRs. The hardness image runs from −1 to +1 and the color coding of the image is so defined. Note in the image that there is a hard, nearly circular edge to the SNR while the interior is considerably softer. The bright central spot does not stand out as being very hard or very soft. The color coding at that location corresponds to ≈0.3 in HR, suggesting that if there is a point source present, it is not a typical hard pulsar like the Crab.

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Figure 3 shows radial profiles of the Chandra image, quantifying the shape of the shell. Given the circular nature of S8, the radial profiles were extracted from a center determined by the mean circular outer edge. The surface brightnesses along the southeast, northeast, southwest, and northwest radii were then plotted. The choice of direction for the radial cuts was made to test the elliptical shape reported by Lozinskaya et al. (1998) for the HRI and Hα images. There are differences in the profiles, but most of the differences are confined to the interior.

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That the profiles demonstrate approximately equal emission throughout the remnant argues for a plerionic or young composite interpretation of the SNR. Rosado et al. (2001) argued against that interpretation, based on the available ROSAT data as well as optical and radio images. They suggested that S8 represented only part of a shell of an older remnant, with much of the SNR hidden by dust. Given Chandra's spatial resolution and the observation presented in this paper, we do not agree with Rosado et al. (2001).

3.3. Radio and Hα Images

S8 has been observed in the radio and Hα bands and it is resolved in both bands (Rosado et al. 2001; Lozinskaya et al. 2003; Lozinskaya & Podorvanyuk 2008). Figure 4 shows contours from the Chandra image overlaid on an Hα data cube described by Lozinskaya et al. (2003). While the peak of the radio emission consistently lies east of the X-ray peak, it does not appear to be spatially distinct or resolved. The Hα data cube of Rosado et al. (2001) is very similar. Overall, the X-ray and Hα contours are spatially coincident and do not demonstrate significant distortion from a circular form. It is possible that a higher spatial resolution version of Figure 4 could determine the location of a compact object. Certainly that the peaks are approximately in the same location warrants a sensitive, high-resolution observation at that location.

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A comparison of the Chandra and radio data yield an image similar to Figure 4, which is not included here. S8 is much less luminous than other radio SNRs in our Galaxy and nearby irregular galaxies: the radio luminosity of S8 is 5% and 17% of the radio luminosity of Cas A and the Crab, respectively, and 15% of the radio luminosity of LMC N 49, the brightest SNR in the LMC (Becker et al. 1991; Feast 1991; Reed et al. 1995).

The lower radio luminosity of S8 is probably due to the SNR shock expanding into a low-density ISM as shown by the H i measurements of Lozinskaya & Podorvanyuk (2008). The SNR is positioned on the edges of a series of H i arcs as noted in their data.

The spectrum shows considerable emission in the 0.7–1.0 keV band, likely unresolved emission lines of Fe or Ne (Figure 5(a)). There are zero events below ∼0.4 keV or above ∼3–4 keV.

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We carried out spectral fits using XSPEC (Arnaud 1996) using the cstat statistic. Fitted parameter values are listed in Table 1.

Table 1.  Spectral Fits for the SNR^a

									Fluxc	Unfolded Fluxc
				Primary	NH	Other		Model	0.5–2 keV	0.5–2 keV
Model	DoF	cstat	χ2/DoF	Parameterb	1022 cm−2	Parameterb	Abundance	Norma	2–8 keV	2–8 keV
Apec	39	79.4	2.06	T: 0.59(4)	<0.6	⋯	all: <1.10	1.4(6)[−5]	⋯	⋯
v-Apec	37	45.3	1.05	T: 0.43(6)	<0.15	⋯	O: 0.46(26)	5.(9)[−5]	5.2[−14]	5.6[−14]
							Fe: 0.15(6)		1.2[−14]	1.2[−14]
Sedov	38	35.3	0.81	T: ${0.98}_{-0.61}^{+0.31}$	<0.09	Tb: <1.4	all: <1.05	5.4[−5]	5.3[−14]	5.5[−14]
						τ: ${8.5}_{-2.8}^{+3.7}$[11]		1.3[−14]	1.3[−14]

Notes.

^aNumbers in () = symmetric uncertainty in last digit(s); numbers in [] = 10^XX. ^bUnits of T and T_b = keV; units of τ = s cm⁻³. ^cFluxes for best-fit models only; all fluxes in erg s⁻¹ cm⁻².

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With an absorption component included, we tried several different models including power law, thermal bremsstrahlung, optically thin thermal gas models (apec, vapec, where v indicates individually variable abundances), and Sedov shock models (Sedov, vSedov; Borkowski et al. 2001). The absorption component used the Wilms et al. (2000) abundance values.

The brems and power-law fits provided poor descriptions of the spectrum and are not included in Table 1. The power-law index was steep (∼7.7 ± 0.9) and that steepness establishes the spectrum as a thermal one, as expected from an SNR.

The best-fit models were the variable-abundance optically thin hot gas model (Smith et al. 2001; Foster et al. 2010) and the Sedov shock model (Borkowski et al. 2001). We tested individually varying abundances in the vsedov model, but found that all of the abundance differences were consistent with solar. Consequently, we do not include this model in Table 1.

The remaining two models (vapec, Sedov) describe the spectrum equally well, but the fitted parameters differ. One may criticize the use of the variable variant of the vapec model given the number of counts in the spectrum. We argue however that the determination of possible nonsolar abundances outweighs the low signal-to-noise. We do not ignore the low signal-to-noise as we fit or fix the abundance of each element sequentially rather than simultaneously.

Furthermore, the nonvariant apec model delivers a visibly poor fit with clear residuals concentrated near the O lines at 0.6 keV and the Fe/Ne lines at 1 keV. We infer from this that, to obtain a good fit to the spectrum, we need at least one additional parameter beyond varying the temperature and column density—whether varying an abundance (vapec) or an ionization time (Sedov). This then also explains why we do not see significant abundance variations with the varying-abundance vSedov model: much of the power of an added parameter is taken up by the ionization time in the Sedov variants. Consequently, additionally varying abundances does not lead to a result significantly different from solar.

The remainder of the spectral discussion will focus on just these two models. Figure 6(a) shows the contours on N_H and kT for the vapec fit and are representative of similar plots for the Sedov model.

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The fitted column density is an upper limit: the best values are ≲0.9 × 10²¹ cm⁻² (Sedov) and ≲1.5 × 10²¹ cm⁻² (vApec) essentially because of the lack of events below ∼0.4 keV. If we assume that the column density lies immediately below those limits, then using the N_H − E_(B−V) relation of Predehl & Schmitt (1995) leads to an estimate of E_(B−V) ∼ 0.17–0.28. Given that IC 1613 is a low-metallicity galaxy, we can also argue for the adoption of the N_H − E_(B−V) of Fitzpatrick (1985; for the Small Magellanic Cloud, SMC) which leads to E_(B−V) ∼ 0.11–0.17. Those values represent an enhancement in the column by a factor of ≈3–8 above the known Galactic column toward IC 1613.

Rosado et al. (2001) infer that the SNR could be partially hidden by dust, so an enhanced column density should be expected. No values are given in their paper. If, however, we crudely assume a spheroid or cuboid shape for IC 1613 and that S8 lies roughly on the mid-plane, then the column length within IC 1613 is ≈5 kpc based on the tidal radius of Battinelli et al. (2007). With a hydrogen number density of 1 cm⁻³ across 1 kpc yielding a column of ≈3 × 10²¹ cm⁻², the numbers work: either a number density of a few per cm⁻³ over a shorter column or a lower value ($\approx 0.1\,{\mathrm{cm}}^{-3}$) over a longer column. Given the appearance of the region surrounding S8 as well as the lower H i density in the core of IC 1613 compared to the northeast complex (Lake & Skillman 1989; Berger et al. 2018), the first case is the more likely, with dust in the region of the northeast complex.

The fitted temperatures are ∼0.43 keV for the vapec model and ∼0.98 keV for the Sedov model. The uncertainty on the vapec model is ≈15% versus a significantly larger uncertainty of ≈30%–80% for the Sedov model. Lacking a basis to choose one over the other, we will explore the implications of temperatures in the 0.4–1 keV range.

We also checked for altered abundances. The simple Sedov model provides an upper limit that is consistent with solar, while the variable vapec model exhibits two nonsolar abundances. Figure 6(b) shows the contours for the vapec model for the two elements with significant nonsolar abundances: O and Fe. The oxygen value is ≈0.5 with an uncertainty of ∼60% while the Fe value is ≈0.15 with an uncertainty of ∼40%.

The Sedov model has an extra parameter, the ionization time, the product of the SNR's age and its post-shock electron number density. Both the simple and the variable-abundance versions returned fitted values of the ionization time, $\sim 8.5\mbox{--}9.5\times {10}^{11}{\rm{s}}\,{\mathrm{cm}}^{-3}$. This indicates that the SNR is close to collisional ionization equilibrium (CIE), validated by the acceptable fit with the vapec model, which itself assumes CIE. That value fits within the picture described in Smith & Hughes (2010).

The integrated flux of the SNR in the 0.3–8 keV band for the best-fit models is ≈5.2–5.3 × 10⁻¹⁴ erg s⁻¹ cm⁻²; the unfolded flux in that band is 5.4–5.6 × 10⁻¹⁴ erg s⁻¹ cm⁻² corresponding to an unfolded luminosity of ≈5.6 × 10³⁶ erg s⁻¹. Our integrated L_X is 25% higher than the estimate in Lozinskaya et al. (1998). In that paper, the flux was estimated in the 0.1–2.4 keV for the ROSAT HRI, an instrument without spectral resolution, using several basic continuum models (e.g., bremsstrahlung, Raymond–Smith, power law). As we have shown, there is little flux below ∼0.4 keV or above ∼2.5 keV, so the two estimates should be considered to be identical.

Shock physics allows us to take the numbers we have measured and turn them into estimates of the explosion energy and age of the SNR. We use expressions from Lequeux (2005) for the shock radius r_s and the shock temperature T_s:

$$\begin{eqnarray*}\begin{array}{rcl}{r}_{{\rm{s}}} & = & 0.26{\left(\displaystyle \frac{{n}_{{\rm{H}}}}{{\mathrm{cm}}^{-3}}\right)}^{-\tfrac{1}{5}}{\left(\displaystyle \frac{t}{{\rm{yr}}}\right)}^{\tfrac{2}{5}}{\left(\displaystyle \frac{E}{4\times {10}^{50}\mathrm{erg}}\right)}^{\tfrac{1}{5}}\,\mathrm{pc}\\ {T}_{{\rm{s}}} & = & 1.5\times {10}^{11}{\left(\displaystyle \frac{{n}_{{\rm{H}}}}{{\mathrm{cm}}^{-3}}\right)}^{-\tfrac{2}{5}}{\left(\displaystyle \frac{t}{\mathrm{yr}}\right)}^{-\tfrac{6}{5}}{\left(\displaystyle \frac{E}{4\times {10}^{50}\mathrm{erg}}\right)}^{\tfrac{1}{5}}\,{\rm{K}},\end{array}\end{eqnarray*}$$

where n_H represents the number density, t represents age, and E represents the explosion energy. We have measures for r_s and T_s; we need an estimate for the number density to proceed. The model normalization provides one estimate for n_H; another estimate comes from the optical [S ii]/Hα ratio.

The model normalizations are scaled identically with the normalization defined equal to

$$\begin{eqnarray*}&&\displaystyle \frac{{10}^{-14}}{4\pi {[{D}_{{\rm{A}}}(1+z)]}^{2}}\int {n}_{{\rm{e}}}{n}_{{\rm{H}}}{dV},\end{eqnarray*}$$

where D_A is the angular diameter distance, n_e = 1.2n_H is the number densities of electrons and protons, respectively, V is the emission volume, and z is the redshift, which is  0 for IC 1613. Inserting numbers leads to an estimate of n_H ∼ 1.6 ± 0.2 cm⁻³ for the vapec and Sedov models.

With the number density estimates, we are then left with two equations with two unknowns: the age t and the explosion energy E. Inserting all numbers leads to

$$\begin{eqnarray*}&&t\sim 3380\mbox{--}5650\,\mathrm{yr}\end{eqnarray*}$$

and

$$\begin{eqnarray*}&&E\sim 3.5\mbox{--}9.9\times {10}^{51}\,\mathrm{erg}\end{eqnarray*}$$

for the age and explosion energy, respectively, and where the range results from the range in temperature (0.4–0.98 keV). Age and explosion energy trade off here, so if we use the estimate of the optical number density, the age range shrinks to 1440–2410 yr and the explosion energy becomes ridiculously high, with a lower limit of 10⁵⁴ erg. This is clearly nonsensical. We therefore conclude that the number density in the X-ray-emitting region is low with n_H ∼ 1.5–2 cm⁻³ confirming the results of Lozinskaya et al. (1998).

With those numbers in hand, the shock velocity is then ∼660 km s⁻¹ for the vapec model or ∼1100 km s⁻¹ if we use the Sedov model numbers. Based on the analysis of Lozinskaya et al. (1998), the lower value is favored.

In Table 2, several basic properties of S8 are compared with other Local Group SNRs. For similar ages of 3–6 kyr, the number densities generally lie in the ≈0.1–1 cm⁻³ range. The luminosity of S8 is the least luminous of the SNRs in the table, likely reflecting the lower interstellar medium density in a dwarf galaxy.

Table 2.  Comparison of SNR S8 with Other Local Group SNRs^a

SNR	Ts	NH	Lx	Ds	nH	Age	E0
S8b	0.43	<1.5	3.5	19	1.7	3.8	2
	0.98	<0.9	⋯	⋯	1.8	6.5	7
NGC 6822c	${2.8}_{-2.0}^{+6.1}$	${3.0}_{-0.7}^{+1.9}$	8.1	24	${0.03}_{-0.02}^{+0.06}$	${3.02}_{-1.32}^{+2.74}$	0.3AD
M33 SNR21d	0.46(2)	0.5 + <0.3d	15	10	1.7	6.7	1.8
LMC N63Ae	0.6	1.7 ± 0.1	⋯	16.4	5	fewe	⋯

Notes.

^aUnits: T in keV; N_H = absorption column in units of 10²¹ cm⁻²; L_x units = 10³⁶ erg s⁻¹ in the 0.5–8 keV band; D_s = SNR diameter in pc; n_H = number density in cm⁻³; Age in kyrs; E₀ = explosion energy in units of 10⁵¹ erg; AD = ADopted value. ^bThis paper. ^cFrom Kong et al. (2004). ^dFrom Gaetz et al. (2007) using the entire SNR field 4 data set for the sedov model. The N_H values represent the Galactic column + the local column. ^eFrom Chu et al. (1999) and Warren et al. (2003).

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We can assign an upper limit on the presence of a pulsar at the center of S8. We adopt two different types of central object: a Crab-like pulsar and a Cas A-like compact object. Assuming that we may treat a compact object as filling a single Chandra pixel and because we do not know the emission properties in the center of the SNR, the detected flux in the center constrains our upper limit. The counts in those two pixels are 35 ± 7.0, for a count rate of ∼7.1 ± 1.4 × 10⁻⁴ cts s⁻¹ and 53 ± 8.3 for a count rate of ∼1.1 ± 0.2 × 10⁻³ cts s⁻¹. Both values are well above the interpolated background counts at the pixels' locations and above the limiting sensitivity values for a point source.

Adopting a soft spectrum appropriate for a Cas A-like object leads to an upper flux limit of ∼5.7 × 10⁻¹⁵ erg s⁻¹ cm⁻² or a 0.3–5 keV luminosity of ∼3.6 × 10³⁵ erg s⁻¹. A Cas A-like object has a luminosity in a similar band of ∼few × 10³³ erg s⁻¹ (Chakrabarty et al. 2001). On the basis of this estimate, we could not detect a Cas A-like object, failing by ≈2 orders of magnitude. Using those same numbers, however, we would easily detect a Crab-like pulsar with a typical luminosity 10–100 times above our sensitivity value. If we estimate the flux from a harder spectrum more typical of a Crab-like pulsar, then our limiting luminosity rises slightly to ∼9 × 10³⁵ erg s⁻¹, but still well below the actual luminosities of Crab-like pulsars, so, again, easily detected.

We note that S8 is detected in the X-ray, optical, and radio bands and behaves rather similarly in all of them. This behavior is counter to the conclusion of Pannuti et al. (2007) in which detections of SNRs in nearby galaxies in all three bands form a very sparse set, in contrast to the number of SNRs detected in any one of those bands. Without additional data, the explanation could always be attributed to insufficient sensitivity of the detectors. However, the observations presented here of S8 at least partially demonstrate that the necessary sensitivity is present, again raising the puzzle described by Pannuti et al. (2007).

Given the X-ray, optical, and radio detections, what type of SNR is S8? The circular, filled X-ray image argues for a plerion or young composite SNR. The spectral indices of the two radio observations, while consistent with each other, are not consistent with a plerion which has typical values in the 0.05–0.30 range (for 3C 58, Bietenholz et al. 2001; and the Crab, Bietenholz & Kronberg 1991; Baars et al. 1977) compared to S8's ≈−0.6 to  −0.7 index value. We then infer that S8 is a young composite SNR.

Do irregular or dwarf galaxies affect the evolving SNRs in a manner different from spirals? Bozzetto et al. (2017) displayed Venn diagrams for the irregular galaxies SMC, LMC, NGC 4449, NGC 3077, NGC 4214, and NGC 5204. All except the LMC and the SMC showed minimal overlap in radio, optical, and X-ray bands. The answer to the question would then appear to be "no," which means that the original puzzle remains. Clearly more work and higher sensitivity surveys are needed to understand the behaviors of SNRs across the electromagnetic spectrum.

Ou et al. (2018) examined LMC SNRs on the basis of their size and X-ray luminosity. They showed that small SNRs all had Type Ia progenitors while medium-sized SNRs were dominated by Type II progenitors because the massive stars had evacuated the immediate vicinity of space. On the basis of its size and L_X, S8 falls in a group of five remnants (N132D, N63A, 0540–69.3, N49, and DEM L71), only one of which had a Type Ia progenitor. On the basis of the its appearance, and its size and luminosity, the position of S8 in the Ou et al. (2018) plot suggests it had a massive progenitor.

We have described the Chandra observation of S8, the only known SNR in the irregular Local Group galaxy IC 1613. S8 is visible in the X-ray, Hα, and radio bands. The SNR is resolved with Chandra into a roughly circular structure with a diameter of ∼19 pc and an estimated age of ∼3400–5600 yr. It exhibits enhanced X-ray emission south of the center of the SNR as defined by the roughly circular edge, but the enhancement is relatively soft. There may be a source at or near the center based on the enhanced emission in the region. Folding together the radio spectral index, the circular outer boundary in the X-ray image, and the apparently filled appearance in the center, we suggest that S8 is a young composite SNR.

We thank the referee for comments that improved this paper. This research has made use of NASA's Astrophysics Data System as well as the NASA/IPAC Extragalactic Database (NED), which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. This research has also made use of data obtained from the High Energy Astrophysics Science Archive Research Center (HEASARC), provided by NASA's Goddard Space Flight Center. This work was supported by Chandra Grant GO3-4104Z.

Facility: CXO(ACIS). -