Element Stratification in the Middle-aged SN Ia Remnant G344.7–0.1

3.1.1. Extended Emission

Figure 2(a) shows a three-color, exposure-corrected image of G344.7–0.1, where red, green, and blue correspond to emission from the Si K (1.76–1.94 keV), hard continuum (3.0–6.0 keV), and Fe K (6.35–6.55 keV) band, respectively. The image indicates that the softer X-ray emission (red) spreads across the entire SNR, whereas the hard X-ray continuum (green) is relatively faint in the northeast region, implying a temperature gradient from the southwest (high) to the northeast (low). This gradient is more clearly highlighted in Figure 2(b), a map of the mean energy of the detected photons in the broad energy band of 1.5–5.0 keV generated using the mean_energy_map script in the CIAO package.¹² Note, however, that this energy band includes several emission lines, and thus a more quantitative spectral analysis (Section 3.2) is required to determine the accurate temperature distribution.

Zoom In Zoom Out Reset image size

Figures 2(c) and (d) are the narrowband images of the Si K and Fe K emission, identical to the red and blue images in Figure 2(a), respectively. It is remarkable that the Fe K line-enhanced region is surrounded by the arc-like structure of the Si K emission. Also notable is that the Fe K peak location is about 2' offset to the west with respect to the geometric center of the SNR determined by the radio band image. In addition to the arc of the Si K emission, a filamentary structure is found at the west rim in the same energy band, indicated with the dashed ellipse in Figure 2(c). All of these features are discovered in this work for the first time, owing to the deep observations with the Chandra ACIS-I.

We generate in Figure 3 radial profiles of the surface brightness in the narrow energy bands that correspond to the K-shell emission of Si (1.76–1.94 keV), S (2.25–2.60 keV), Ar (2.95–3.25 keV), Ca (3.60–4.10 keV), and Fe (6.35–6.55 keV). We confirm that the Fe K emission profile indeed peaks at the interior to other species. We also find that the profiles of the intermediate-mass elements (IMEs) are similar to each other.

Zoom In Zoom Out Reset image size

3.1.2. Point-like Sources

Figure 2(e) shows a 0.5–1.2 keV image of the same region as the other panels of Figure 2. Because of the large foreground extinction (N_H ∼ 5 × 10²² cm⁻²: Yamauchi et al. 2005; Combi et al. 2010; Giacani et al. 2011; Yamaguchi et al. 2012), soft X-rays from G344.7–0.1 are almost fully absorbed. Therefore, photons in this energy band are likely dominated by foreground objects. Our simple eyeball inspection detects the two bright point-like sources, labeled No. 1 and 2 in Figure 2(e). Source No. 1 corresponds to CXOU J170357.8–414302, previously reported by Combi et al. (2010). This source is located near the geometric center of the SNR, and does not positionally coincide with the Fe K peak location. The other (No. 2) is newly detected with our deep Chandra observation. We also search for fainter point sources using the wavdetect algorithm in the two energy ranges of 0.5–1.2 keV and 1.2–2.5 keV with the Encircled Counts Fraction (ECF) parameter set to 0.9. Table 2 gives the detected sources with their coordinates, spatial extent relative to the size of the point-spread function (PSFRATIO), and the counts in the 0.5–2.0 keV (soft) and 2.0–9.0 keV (hard) bands. We regard those with PSFRATIO <1 as true point sources. About 50 sources are detected based on these criteria.

Table 2.  List of the Detected Point-like Sources

No.	${\left({\rm{R}}.{\rm{A}}.,\mathrm{decl}.\right)}_{{\rm{J}}2000}$	PSFRATIOa	Softb	Hardc	No.	${\left({\rm{R}}.{\rm{A}}.,\mathrm{decl}.\right)}_{{\rm{J}}2000}$	PSFRATIOa	Softb	Hardc
1	17:03:57.8, −41:43:02.2	0.56	146	44	26	17:03:36.8, −41:40:46.5	0.35	29	65
2	17:04:04.3, −41:41:01.0	0.48	114	38	27	17:03:54.4, −41:40:19.9	0.70	31	59
3	17:03:47.5, −41:46:47.1	0.77	65	67	28	17:04:09.3, −41:40:11.3	0.69	22	29
4	17:04:04.1, −41:45:36.2	0.24	24	2	29	17:04:19.3, −41:38:56.3	0.84	29	85
5	17:04:01.2, −41:43:35.6	0.25	8	2	30	17:03:55.1, −41:46:33.7	0.90	56	110
6	17:03:53.4, −41:42:32.2	0.34	35	27	31	17:03:53.9, −41:41:33.3	0.62	36	80
7	17:03:45.7, −41:41:56.4	0.43	46	33	32	17:03:49.8, −41:41:01.8	0.81	78	128
8	17:04:09.1, −41:39:56.9	0.64	28	31	33	17:03:38.5, −41:45:15.2	0.64	41	124
9	17:04:07.0, −41:39:49.6	0.72	32	17	34	17:03:37.2, −41:44:36.8	0.80	137	306
10	17:04:11.8, −41:39:31.0	0.33	17	1	35	17:03:52.7, −41:44:27.8	0.51	20	32
11	17:03:41.6, −41:46:28.9	0.47	21	39	36	17:03:35.9, −41:44:22.0	0.51	49	118
12	17:04:03.0, −41:44:06.2	0.28	8	3	37	17:03:46.8, −41:44:13.8	0.68	39	109
13	17:04:11.7, −41:39:53.4	0.28	11	4	38	17:03:44.7, −41:43:57.6	0.37	15	34
14	17:04:20.1, −41:39:47.7	0.35	28	13	39	17:03:34.4, −41:43:20.2	0.67	80	170
15	17:04:18.5, −41:39:21.0	0.38	46	22	40	17:03:55.2, −41:43:12.4	0.30	7	7
16	17:04:16.3, −41:38:37.9	0.36	24	7	41	17:03:56.1, −41:43:00.6	0.38	10	20
17	17:04:14.4, −41:38:34.2	0.62	80	70	42	17:03:43.9, −41:42:43.3	0.48	25	57
18	17:03:38.8, −41:46:36.1	0.45	40	38	43	17:03:37.0, −41:42:20.0	0.76	116	191
19	17:03:36.8, −41:39:40.1	0.55	35	44	44	17:04:00.3, −41:42:07.1	0.51	8	14
20	17:03:51.7, −41:44:45.8	0.95	118	187	45	17:03:46.1, −41:40:52.3	0.24	14	20
21	17:04:04.7, −41:43:47.9	0.39	15	9	46	17:03:43.4, −41:40:55.5	0.60	70	130
22	17:04:02.3, −41:46:23.0	0.60	39	39	47	17:03:50.4, −41:40:16.8	0.90	105	160
23	17:04:19.7, −41:42:52.4	0.43	19	10	48	17:04:11.1, −41:40:01.8	0.91	51	69
24	17:03:33.9, −41:45:49.8	0.60	68	108	49	17:03:53.4, −41:38:37.4	0.91	41	130
25	17:04:14.0, −41:41:28.7	0.79	22	31	⋯	⋯	⋯	⋯	⋯

Notes.

^aA parameter of wavdetect that indicates the spatial extension of the sources with respect to the PSF size. The source can be regarded as a true point source when this value is unity or less. ^bPhoton counts in 0.5–2.0 keV. ^cPhoton counts in 2.0–9.0 keV.

Download table as:  ASCIITypeset image

3.2.1. Extended Emission

Here we analyze ACIS-I spectra extracted from the regions given in Figure 2, where the point sources detected in the previous section (Table 2) are excluded. In our analysis we model the background emission instead of subtracting it. Figure 4 shows the spectra of several representative regions, where black and red are the data from the May and July observations, respectively. The emission lines of the IMEs and Fe are clearly resolved in each region. Background data are extracted from an annulus region surrounding the SNR with inner and outer radii of $5\buildrel{\,\prime}\over{.} 0$ and $7\buildrel{\,\prime}\over{.} 0$, respectively (Figure 1). We model and fit the background spectra simultaneously with the source spectra, instead of subtracting them from the source data, to properly take into account the difference in the effective area between the source and background regions. The details of the background spectral modeling are presented in Appendix A. In the following spectral analysis, we fit unbinned data using the C-statistics (Cash 1979) to estimate the model parameters and their error ranges without bias, although the binned spectra are shown in the figures for clarity.

Zoom In Zoom Out Reset image size

We first fit the spectra with a model of an optically thin thermal plasma in the non-equilibrium ionization (NEI) state, i.e., the vnei model in the XSPEC package, based on the latest atomic database (AtomDB¹³ version 3.0.9). We use the tbnew_gas model¹⁴ for the foreground absorption, assuming the photoelectric absorption cross sections taken from Verner et al. (1996). The free parameters are the absorption column density N_H, electron temperature kT_e, ionization timescale n_et, volume emission measure (VEM), and elemental abundances of Mg, Si, S, Ar, Ca, and Fe relative to the solar abundance table of Wilms et al. (2000). This model gives good fits to the spectra below 5 keV, but fails to reproduce the Fe K emission. We find that the model generally predicts an Fe K centroid energy higher than the observed values, which indicates that the Fe K emission originates from another plasma component with a lower ionization state.

We thus add a Gaussian component to reproduce the K-shell emission from the low-ionized Fe. The addition of this component significantly improves the fit, with the C-stat value reduced from 1003.00/805 to 920.43/803 in the 5.0–8.0 keV band, including the strong Fe K line. The best-fit model spectra and parameters are given in Figure 4 and Table 3, respectively. The abundances of the IMEs (but for Mg) are generally higher than the solar values, implying a significant ejecta contribution to the vnei component. We also find that both the electron temperature (of the vnei component) and the Fe K centroid energy vary among the regions. The lowest temperature is observed in the N and NE, consistent with the mean photon energy map (Figure 2(b)). The Fe K centroid energy is marginally lower in the Fe peak region than in most of the outer regions, indicating that the Fe ejecta in the former has a low ionization degree. The absorption column density is found to be almost uniform throughout the SNR, ∼6.5 × 10²² cm⁻². This value is slightly higher than the previous measurements of N_H = 4.0–5.5 × 10²² cm⁻² (Yamauchi et al. 2005; Combi et al. 2010; Giacani et al. 2011; Yamaguchi et al. 2012). This discrepancy is likely due to the difference in the reference solar abundance; the previous work referred to Anders & Grevesse (1989), while we refer to Wilms et al. (2000). In fact, if we fit the ACIS-I spectra using the solar abundance table of Anders & Grevesse (1989), we obtain N_H ∼ 4.5 × 10²² cm⁻², consistent with the previous work.

Table 3.  Best-fit Spectral Parameters for the Extended Emission

Components	Parameters	Fe	Si N	Si E	Si SW	N
Absorptiona	${N}_{{\rm{H}}}\ ({10}^{22}\,{\mathrm{cm}}^{-2})$	${6.40}_{-0.24}^{+0.34}$	${6.68}_{-0.29}^{+0.29}$	${6.47}_{-0.20}^{+0.22}$	${6.52}_{-0.17}^{+0.18}$	${6.28}_{-0.50}^{+0.48}$
NEI plasmab	${{kT}}_{{\rm{e}}}\ (\mathrm{keV})$	${1.10}_{-0.08}^{+0.04}$	${0.83}_{-0.04}^{+0.05}$	${0.93}_{-0.03}^{+0.03}$	${0.99}_{-0.03}^{+0.03}$	${0.81}_{-0.07}^{+0.08}$
	Mg	${0.45}_{-0.24}^{+0.33}$	${1.14}_{-0.31}^{+0.40}$	${0.33}_{-0.19}^{+0.23}$	${0.63}_{-0.18}^{+0.21}$	${1.11}_{-0.60}^{+1.16}$
	Si	${1.51}_{-0.16}^{+0.15}$	${1.42}_{-0.12}^{+0.14}$	${1.13}_{-0.94}^{+0.10}$	${1.48}_{-0.09}^{+0.10}$	${1.87}_{-0.17}^{+0.45}$
	S	${2.28}_{-0.17}^{+0.15}$	${1.69}_{-0.11}^{+0.11}$	${1.49}_{-0.09}^{+0.09}$	${2.03}_{-0.08}^{+0.09}$	${2.13}_{-0.23}^{+0.27}$
	Ar	${2.05}_{-0.24}^{+0.26}$	${2.18}_{-0.27}^{+0.29}$	${1.62}_{-0.19}^{+0.20}$	${2.14}_{-0.16}^{+0.16}$	${1.71}_{-0.49}^{+0.54}$
	Ca	${4.26}_{-0.54}^{+0.58}$	${4.09}_{-0.96}^{+1.22}$	${1.85}_{-0.41}^{+0.44}$	${4.41}_{-0.47}^{+0.49}$	${1.82}_{-1.61}^{+2.05}$
	${n}_{{\rm{e}}}t\ ({10}^{11}\ {\rm{s}}\ {\mathrm{cm}}^{-3})$	${1.74}_{-0.36}^{+0.49}$	${1.48}_{-0.31}^{+0.37}$	${2.17}_{-0.38}^{+0.45}$	${1.49}_{-0.16}^{+0.21}$	${1.67}_{-0.63}^{+0.85}$
	VEMc	${8.22}_{-0.84}^{+1.87}$	${18.9}_{-2.9}^{+3.2}$	${17.5}_{-1.7}^{+1.8}$	${26.5}_{-2.4}^{+2.4}$	${5.90}_{-1.40}^{+1.96}$
Fe K	E (keV)	${6.41}_{-0.03}^{+0.03}$	${6.48}_{-0.06}^{+0.06}$	${6.45}_{-0.03}^{+0.03}$	${6.40}_{-0.03}^{+0.03}$	${6.54}_{-0.07}^{+0.07}$
	σしぐま (eV)	${206}_{-31}^{+35}$	100 (fix)	<93	${85}_{-40}^{+39}$	100 (fix)
	Fluxd	${9.80}_{-1.1}^{+1.1}$	${1.65}_{-0.62}^{+0.65}$	${3.53}_{-0.66}^{+0.70}$	${6.77}_{-1.07}^{+1.16}$	${1.02}_{-0.68}^{+0.75}$
Photon Counts (104 conut)		1.75	1.83	2.05	4.22	1.24
Background Contribution (%)		18.4	19.1	16.4	21.4	48.8
C-stat		2658.68	2690.62	2636.78	2739.07	2655.14
dof		2311	2312	2311	2311	2312
		NE	NW	E	SE	Filament
Absorptiona	${N}_{{\rm{H}}}\ ({10}^{22}\,{\mathrm{cm}}^{-2})$	${6.19}_{-0.29}^{+0.31}$	${5.30}_{-0.54}^{+0.75}$	${6.63}_{-0.43}^{+0.60}$	${6.69}_{-0.53}^{+0.57}$	${5.80}_{-0.47}^{+0.84}$
NEI plasmab	${{kT}}_{{\rm{e}}}\ (\mathrm{keV})$	${0.79}_{-0.05}^{+0.07}$	${1.27}_{-0.25}^{+0.25}$	${0.88}_{-0.10}^{+0.06}$	${0.90}_{-0.07}^{+0.07}$	${1.41}_{-0.21}^{+0.24}$
	Mg	${1.25}_{-0.36}^{+0.49}$	<1.70	${1.18}_{-0.55}^{+0.81}$	${0.66}_{-0.47}^{+0.70}$	<1.23
	Si	${2.09}_{-0.22}^{+0.27}$	${3.03}_{-0.76}^{+0.89}$	${1.87}_{-0.28}^{+0.34}$	${1.61}_{-0.25}^{+0.31}$	${2.22}_{-0.47}^{+0.73}$
	S	${2.00}_{-0.15}^{+0.17}$	${4.07}_{-0.80}^{+0.86}$	${2.49}_{-0.25}^{+0.27}$	${2.60}_{-0.24}^{+0.28}$	${2.67}_{-0.43}^{+0.53}$
	Ar	${2.83}_{-0.40}^{+0.46}$	${3.72}_{-0.90}^{+1.03}$	${2.74}_{-0.47}^{+0.51}$	${3.16}_{-0.42}^{+0.48}$	${2.15}_{-0.64}^{+0.73}$
	Ca	${5.73}_{-1.61}^{+2.04}$	${3.28}_{-1.30}^{+1.52}$	${5.22}_{-0.84}^{+1.07}$ e	${6.65}_{-1.58}^{+2.61}$	${4.39}_{-1.46}^{+1.82}$
	${n}_{{\rm{e}}}t\ ({10}^{11}\ {\rm{s}}\ {\mathrm{cm}}^{-3})$	${1.52}_{-0.35}^{+0.41}$	${1.59}_{-0.47}^{+1.21}$	${1.26}_{-0.33}^{+0.79}$	${1.46}_{-0.43}^{+0.57}$	${0.86}_{-0.23}^{+0.39}$
	VEMc	${11.1}_{-2.2}^{+2.3}$	${1.09}_{-0.35}^{+0.75}$	${7.04}_{-1.28}^{+3.05}$	${8.00}_{-1.76}^{+2.19}$	${0.91}_{-0.25}^{+0.41}$
Fe K	E (keV)	${6.45}_{-0.03}^{+0.03}$	⋯	${6.46}_{-0.03}^{+0.03}$ e	${6.46}_{-0.03}^{+0.03}$ e	⋯
	σしぐま (eV)	${93.9}_{-35}^{+36}$	⋯	${106}_{-36}^{+35}$	${167}_{-58}^{+71}$	⋯
	Fluxd	${5.22}_{-0.86}^{+0.94}$	⋯	${5.31}_{-1.01}^{+1.10}$	${3.88}_{-1.00}^{+1.10}$	⋯
Photon Counts (104 conut)		3.73	0.69	1.48	1.52	0.40
Background Contribution (%)		35.5	45.6	46.5	40.6	35.8
C-stat		2704.56	2694.10	2601.55	2643.46	2671.64
dof		2311	2314	2313	2312	2314

Notes.

^aThe absorption cross sections are taken from Verner et al. (1996). ^bSolar abundances of Wilms et al. (2000) are assumed. ^cIn units of ${10}^{11}\ {\mathrm{cm}}^{-5}$. The volume emission measure (VEM) is given as $\int {n}_{{\rm{e}}}{n}_{{\rm{H}}}{dV}/(4\pi {D}^{2})$, where V and D are the volume of the emission region (cm³) and the distance to the emitting source (cm), respectively. ^dIn units of ${10}^{-6}\ \mathrm{photon}\,{\mathrm{cm}}^{-2}\,{\mathrm{keV}}^{-1}\ {{\rm{s}}}^{-1}$. ^eThese are determined using spectra extracted from a larger region containing both E and SE, since the values cannot be constrained for each individual region.

Download table as:  ASCIITypeset image

Next, we replace the Gaussian component with another vnei component to account for the emission from pure Fe ejecta. The SNR is now divided into three regions, the Fe peak, Si-arc (containing Si E, N, and SW), and rim (containing N, NE, NW, SE, S, and Filament), and their spectra are jointly fitted with a common N_H value. Table 4 gives the parameter values obtained for the pure Fe ejecta component. We find that the temperature of this component (kT_e > 2.8keV) is significantly hotter than that of the IME component (kT_e = 0.8–1.1 keV). We also find that n_et of the Fe ejecta <2 × 10¹⁰ s cm⁻³ is lower than that of the IME component ∼2 × 10¹¹ s cm⁻³, and that the lowest n_et of the Fe ejecta component is achieved at the Fe peak region.

Table 4.  Best-fit Spectral Parameters for the Pure Fe Component

Parameters	Fe Peak	Si-arc	Rim
${{kT}}_{{\rm{e}}}\ (\mathrm{keV})$	${2.9}_{-0.1}^{+0.1}$	${5.6}_{-0.4}^{+5.6}$	${3.8}_{-0.2}^{+0.2}$
${n}_{{\rm{e}}}t\ ({10}^{9}\ {\rm{s}}\ {\mathrm{cm}}^{-3})$	<0.14	${3.5}_{-1.8}^{+1.8}$	${17.4}_{-2.2}^{+2.1}$
VEMa $({10}^{5}\ {\mathrm{cm}}^{-5})$	${1.2}_{-0.1}^{+0.1}$	${0.8}_{-0.1}^{+0.4}$	${2.3}_{-0.5}^{+0.1}$

Note.

^aThe VEM is given as $\int {n}_{{\rm{e}}}{n}_{\mathrm{Fe}}{dV}/(4\pi {D}^{2})$.

Download table as:  ASCIITypeset image

3.2.2. Point-like Sources

We perform spectral analysis of two bright sources (No. 1 and No. 2) detected in the soft X-ray band, both embedded in the diffuse emission from the SNR (Figure 2(e)). The spectra of source No. 1 (CXOU J170357.8–414302) and No. 2 are shown in Figure 5. We fit the spectra of both sources with the following three models: an absorbed power law (tbnew_gas×powerlaw), an absorbed blackbody (tbnew_gas×bbody), and an absorbed optically thin thermal plasma with solar abundances (tbnew_gas×apec). We use background spectra extracted from the annulus surrounding each source region. For source No. 1, we obtain 160 counts for the background-corrected spectrum with 16% background contribution, and 136 counts with 8.6% for No. 2. The best-fit parameters we obtain are given in Table 5. The estimated column densities are significantly lower than that of the SNR (${N}_{{\rm{H}}}$ ≈ 6.5 × 10²² cm⁻²) for any model assumption, consistent with the foreground stellar object explanation as suggested by Combi et al. (2010) for No. 1. The spectra of the other detected sources (Table 2) with sufficient photon counts are analyzed as well, obtaining statistically acceptable results (C-stat/dof ≲ 1.2) for all the sources. The results are discussed in more detail in Section 4.3.

Zoom In Zoom Out Reset image size

Table 5.  Best-fit Spectral Parameters for the Bright Point Sources

No.	${N}_{{\rm{H}}}$ a	Γがんま	${N}_{{\rm{H}}}$ a	${{kT}}_{\mathrm{bb}}$ b	${N}_{{\rm{H}}}$ a	${{kT}}_{\mathrm{APEC}}$ b
1	<0.08	${3.44}_{-0.24}^{+0.29}$	<0.06	${0.24}_{-0.01}^{+0.02}$	${0.58}_{-0.18}^{+0.18}$	${0.92}_{-0.14}^{+0.10}$
2	${0.93}_{-0.31}^{+0.35}$	${3.84}_{-0.47}^{+0.53}$	<0.16	${0.43}_{-0.03}^{+0.03}$	${1.29}_{-0.50}^{+0.25}$	${1.04}_{-0.15}^{+0.31}$

Notes.

^aIn units of ${10}^{22}\,{\mathrm{cm}}^{-2}$. ^bIn units of keV.

Download table as:  ASCIITypeset image

Thanks to deep observations with Chandra, we have revealed the centrally peaked Fe K emission surrounded by the arc-like structure of the IME emission. This radially stratified chemical composition with Fe at the interior is consistent with the standard picture of SN Ia nucleosynthesis dominated by carbon detonation (e.g., Iwamoto et al. 1999; Seitenzahl et al. 2013). Several other SN Ia explosion models, such as those involving pure carbon deflagration (e.g., Fink et al. 2014) or He-shell detonation (e.g., Sim et al. 2012), predict a substantial amount of ⁵⁶Ni (which decays into Fe) synthesized outside of the IME. We do not find such evidence in the X-ray data of G344.7–0.1, which may rule out these explosion scenarios as the origin of this SNR. We note that our spectral data below 1 keV, where Fe L lines fall, are missing due to the strong absorption, leaving the possibility that a significant amount of Fe is mixing into the IME-rich layer. Unless this is the case, our results prefer the carbon detonation model for the origin of G344.7–0.1.

We have also found that the peak location of the Fe emission (and thus the centroid of the IME arc) is offset to the west by $\sim 1\buildrel{\,\prime}\over{.} 6$ (corresponding to ∼2.8 pc at the distance of 6 kpc) with respect to the geometric center of the SNR. This implies that either the SN explosion was asymmetric, or the Fe K peak location is the actual explosion center but an apparent asymmetry has formed during the SNR evolution. The spatial distribution of the plasma conditions may help disentangle this degeneracy. In particular, the ionization state of the Fe ejecta that is inferred from the Fe K centroid energy offers a useful diagnostic of the time duration since the ejecta were heated by the reverse shock (e.g., Badenes et al. 2006; Yamaguchi et al. 2014b). Figure 6(a) shows the Fe K centroid energy observed in each region as a function of the angular distance from the location of the Fe emission, which reveals a positive correlation between the two quantities. The mean values of the centroid energy in the Fe peak and Si arc regions correspond to the charge states of Fe¹⁰⁺ and Fe¹⁷⁺, respectively. Similarly, Table 4 indicates that a lower ionization timescale is achieved in the Fe peak region than in the outer regions. This suggests that the Fe ejecta at the inner regions were heated by the reverse shock more recently. We also plot in Figure 6(b) the electron temperatures of the NEI component of each region. In contrast to the Fe K centroid, the temperature is anticorrelated with the distance from the Fe peak location. We thus conclude that the location of the Fe peak region is the true explosion center of this SNR. The origin of the asymmetric morphology is discussed in more detail in Section 4.2.

Zoom In Zoom Out Reset image size

One caveat is that the variation in the Fe K centroids may also be contributed by the line-of-sight velocity of the Fe ejecta, as was claimed for the Kepler SNR (Kasuga et al. 2018). We find that the observed variation in the Fe K centroid (∼50 eV) requires a line-of-sight velocity v_sight of 2400 km s⁻¹ for the Fe ejecta at the Fe peak region. If the Fe ejecta of the Fe peak location have been freely expanding for 3000 yr with v_sight ∼2400 km s⁻¹, the ejecta reach 7.4 pc from the SN explosion center, larger than the radius of this SNR (5 pc; assuming a distance of 6 kpc). Given that the ejecta must have been decelerated by the reverse shock, we should take this estimate as a lower limit. Since an unnaturally large asymmetry in the explosion is required to explain this line shift solely with the Doppler effect, the ionization effect is likely dominant in the observed variation in the Fe K centroid energies. The degeneracy between the ionization effect and Doppler shift in the Fe K emission can be disentangled by measurement of the Fe Kβべーた/Kαあるふぁ flux ratio that strongly depends on the dominant charge state and thus constrains the rest-frame Fe K centroid (Yamaguchi et al. 2014a). Unfortunately, neither our Chandra observation nor any previous work allows us to detect Fe Kβべーた emission due to poor photon statistics and limited energy resolution. Future high-resolution spectroscopy with XRISM and Athena are crucial to determine both charge state and line-of-sight velocity accurately.

Now we estimate the three-dimensional distribution of the vnei component that is dominated by the IME ejecta, assuming that the Fe peak region is the true explosion center. Since the surface-brightness profiles of Si, S, Ar, and Ca (Figure 3) are virtually identical with one another, we use a radial profile of the 1.5–5.0 keV emission as representative of the IME distribution. Figure 7(a) shows the extracted radial profile compared with a simple two-zone spherical model assuming a uniform density and uniform abundance distribution in each zone. We also take into account the temperature gradient, and thus the emissivity variation, between the two zones, where ∼1.1 keV at <1' and ∼0.9 keV at 1'–3'. Figure 7(b) shows the continuum-subtracted surface brightness of the Fe K and 1.5–5.0 keV emission, suggesting an enhanced density of Fe ejecta at <1'. We find that the observed profile is well reproduced by a model consisting of the model schematized in Figure 7(c): a low-density (n_in) inner sphere with a radius of 1' and a thick high-density (n_shell = 2.6n_in) shell with inner and outer radii of 1' and 3', respectively. The IME arc structure found in Figure 2(c) corresponds to the inner part of the thick shell, where the highest surface brightness is achieved. Note that the angular distance between the explosion center and the west rim is ∼3', thus the excess in the brightness beyond this radius (Figure 7(a)) is solely due to the emission extended to the northeast.

Zoom In Zoom Out Reset image size

As discussed in the previous section, the identified explosion center is offset from the geometric center of the SNR. Therefore, the X-ray morphology of G344.7–0.1 is relatively asymmetric. Lopez et al. (2011) classified this SNR as a possible CC SNR because of such highly asymmetric X-ray morphology, based on the analysis of the power-ratio method (PRM). However, the data set used in Lopez et al. (2011) has more limited photon statistics than our 200 ks Chandra data. Thus, we apply the PRM on our new data set of G344.7–0.1 to measure the asymmetry of this SNR. We replicate the analysis from Lopez et al. (2011), where the PRs are calculated on the 0.5–2.1 keV band exposure-corrected image of SNR, and obtain ${P}_{2}/{P}_{0}={9.84}_{-2.32}^{+2.05}$, ${P}_{3}/{P}_{0}={1.75}_{-0.62}^{+0.55}$.¹⁵ The values are now comparable with those for the other well-established Type Ia SNRs, such as Kepler. We note that the surface-brightness-weighted geometric center, i.e., the aperture center for the PRM, is (αあるふぁ_J2000, δでるた_J2000) = (17^h03^m555, −41°42'545), which is almost the same as the geometric center of the SNR, and thus is offset from our proposed explosion center of (αあるふぁ_J2000, δでるた_J2000) = (17^h03^m491, −41°42'460). This offset and apparent asymmetry seems to be caused by the nonuniform distribution of the ambient medium. In fact, other SNRs expanding in such ambient medium distributions (e.g., Kepler, N103B) also show high asymmetric morphology and off-center aperture center like G344.7–0.1 (Lopez et al. 2011).

The nonuniform density structure of the ambient medium is indeed found in our radio data. Figure 8 shows the total interstellar proton column density map ${N}_{{\rm{p}}}({{\rm{H}}}_{2}+{\rm{H}}\,{\rm{I}})$ combining both the molecular (CO: NANTEN2) and atomic (H i: McClure-Griffiths et al. 2005) components. The details of the radio observations and data analysis are described in Appendix B. The velocity range of the gas density map is from −118.0 km s⁻¹ to −109.0 km s⁻¹, which corresponds to the kinematic distance of 6.2–6.4 kpc adopting a model of the Galactic rotation curve by Brand & Blitz (1993). The velocity range is determined by CO/H i cavities in the position–velocity diagrams (see Appendix B.3), suggesting an expanding gas motion originated by accretion winds from the progenitor system of the SNR (e.g., Sano et al. 2017, 2018). Therefore, the molecular and atomic clouds at the velocity range are likely associated with the SNR. The velocity range and distance to the SNR are consistent with the previous H i study by Giacani et al. (2011).

Zoom In Zoom Out Reset image size

We find denser clouds possibly interacting with the SNR in the west (Figure 8), consistent with the previous results from the Mopra Southern Galactic Plane CO Survey (Burton et al. 2013; Braiding et al. 2015; Lau et al. 2019). The SNR expansion is likely to decelerated by this interaction toward the west, resulting in the asymmetric morphology we observe. This interpretation is supported by the temperature gradient inferred from the mean energy map (Figure 2(b)); the lower temperature is achieved in the northeast region probably due to more efficient adiabatic expansion. A similar temperature gradient due to the asymmetric ambient density distribution is observed in other middle-aged SNRs, e.g., W49B (Lopez et al. 2013; Yamaguchi et al. 2018). Alternatively, the SN ejecta at the west could have been reheated by reflection shocks originating from the interaction with the dense clouds, as was observed in, for instance, Kes 27 (Chen et al. 2008).

We have also discovered a local asymmetric structure at the west rim, "Filament" in Figure 2(c). Interestingly, the measured IME abundances are enhanced with respect to the solar values. Moreover, the relative abundances of S, Ar, Ca to Si are comparable to the solar composition and that of Mg is substantially lower, which is consistent with the typical nucleosynthesis yields of the incomplete Si burning (e.g., Iwamoto et al. 1999; Seitenzahl et al. 2013; Townsley et al. 2016). This implies an ejecta origin of this structure, despite its filamentary morphology and location being more consistent with an interpretation of the forward shock front that propagates into the ISM. If the structure is indeed dominated by the SN ejecta, it is reminiscent of the ejecta knots in Tycho SNR, suggesting that such local asymmetry is somewhat common in a certain type of Type Ia SNRs.

We have analyzed the spectra of the point-like sources detected with total photon counts of 30 or more (see Table 2). The best-fit parameters for the absorbed blackbody model are given in Table 6, where the angular distance from the explosion center to the sources is given as well. If the sources are physically associated with the SNR, the measured absorption column density must be consistent with that of the SNR (5–7 × 10²² cm⁻²). We find that only eight sources, No. 31, 36, 39, 41, 42, 43, 48, and 49, satisfy this criterion, and the other sources are most likely either foreground or background objects.

Table 6.  Spectral Properties of All Point Sources

No.a	${N}_{{\rm{H}}}\ ({10}^{22}\,{\mathrm{cm}}^{-2})$	${{kT}}_{\mathrm{bb}}\ (\mathrm{keV})$	Luminosityb (${10}^{33}\ \mathrm{erg}\,{{\rm{s}}}^{-1}$)	Distance $(\mathrm{arcmin})$ c
1	<0.06	${0.24}_{-0.01}^{+0.02}$	${4.85}_{-0.47}^{+0.93}\times {10}^{-2}$	1.66
2	<0.16	${0.43}_{-0.03}^{+0.03}$	${3.53}_{-0.32}^{+0.56}\times {10}^{-2}$	3.34
3	<0.09	${0.58}_{-0.05}^{+0.06}$	${2.58}_{-0.35}^{+0.39}\times {10}^{-2}$	3.99
6	<0.10	${0.59}_{-0.07}^{+0.07}$	${1.02}_{-0.16}^{+0.21}\times {10}^{-2}$	4.69
7	<0.38	${0.28}_{-0.07}^{+0.10}$	${1.46}_{-0.30}^{+1.81}\times {10}^{-2}$	4.45
8	<0.10	${0.84}_{-0.15}^{+0.14}$	${1.11}_{-0.29}^{+0.37}\times {10}^{-2}$	5.35
9	<0.93	${0.16}_{-0.06}^{+0.04}$	${1.86}_{-0.87}^{+12.3}\times {10}^{-2}$	2.93
11	${1596}_{-1008}^{+616}$	<0.24d	$\lt 3.15\,\times \,{10}^{16}$	5.11
14	<0.38	${0.34}_{-0.08}^{+0.07}$	${0.84}_{-0.17}^{+0.62}\times {10}^{-2}$	6.55
15	<0.14	${0.48}_{-0.05}^{+0.05}$	${3.23}_{-0.48}^{+0.61}\times {10}^{-2}$	6.33
16	<0.68	${0.21}_{-0.04}^{+0.05}$	${1.05}_{-0.38}^{+0.80}\times {10}^{-2}$	4.29
17	<0.71	${0.54}_{-0.07}^{+0.08}$	${4.08}_{-0.23}^{+1.01}\times {10}^{-2}$	2.66
18	<2.45	${0.25}_{-0.14}^{+0.10}$	${0.78}_{-0.37}^{+10.6}\times {10}^{-2}$	3.85
19	<0.62	${0.45}_{-0.34}^{+0.17}$	${0.23}_{-0.23}^{+0.24}\times {10}^{-2}$	3.09
20	<0.25	${0.80}_{-0.08}^{+0.07}$	${4.66}_{-0.60}^{+0.69}\times {10}^{-2}$	1.04
22	<0.54	${0.39}_{-0.10}^{+0.10}$	${1.05}_{-0.20}^{+0.81}\times {10}^{-2}$	4.38
24	<1.23	${0.56}_{-0.16}^{+0.11}$	${1.49}_{-0.36}^{+1.27}\times {10}^{-2}$	4.17
25	<0.41	${0.88}_{-0.31}^{+0.44}$	${0.74}_{-0.28}^{+0.44}\times {10}^{-2}$	4.82
26	${2.65}_{-1.20}^{+1.54}$	${0.79}_{-0.13}^{+0.17}$	${4.05}_{-0.90}^{+1.52}\times {10}^{-2}$	3.03
27	${38}_{-4.7}^{+0.8}$	${2.73}_{-0.02}^{+0.14}\times {10}^{-2}$	$\lt 3.60\,\times \,{10}^{29}$	2.63
28	<0.62	${0.48}_{-0.14}^{+0.15}$	${0.51}_{-0.17}^{+0.31}\times {10}^{-2}$	4.58
29	<1.90	>2.26	<253	6.82
30	<0.39	${0.90}_{-0.19}^{+0.27}$	${0.66}_{-0.35}^{+0.38}\times {10}^{-2}$	3.96
31	${7.44}_{-3.29}^{+5.60}$	${0.37}_{-0.11}^{+0.13}$	<2.77	1.51
	6.69e	${0.39}_{-0.04}^{+0.05}$	${0.21}_{-0.08}^{+0.09}$
32	${8.66}_{-3.38}^{+7.79}$	<0.18	$\lt 7.64\,\times \,{10}^{10}$	1.74
33	${38}_{-14}^{+23}$	${0.07}_{-0.02}^{+0.04}$	$\lt 9.00\,\times \,{10}^{11}$	3.17
34	${2.42}_{-0.89}^{+1.19}$	${0.62}_{-0.08}^{+0.09}$	${0.12}_{-0.03}^{+0.05}$	2.89
35	>710	<197	$\lt 5.36\,\times \,{10}^{7}$	1.83
36	${6.74}_{-2.49}^{+3.17}$	${0.35}_{-0.07}^{+0.09}$	<1.12	2.93
	6.63e	${0.35}_{-0.03}^{+0.06}$	${0.24}_{-0.13}^{+0.10}$
37	${2.18}_{-1.29}^{+1.90}$	${0.63}_{-0.14}^{+0.20}$	${3.27}_{-1.18}^{+2.96}\times {10}^{-2}$	1.52
38	${27}_{-13}^{+19}$	${0.12}_{-0.04}^{+0.07}$	$\lt 4.12\,\times \,{10}^{5}$	1.45
39	${3.60}_{-1.78}^{+2.73}$	${0.49}_{-0.12}^{+0.14}$	${5.67}_{-2.97}^{+12.1}\times {10}^{-2}$	2.80
41	<28.3	<4.19	${1.01}_{-0.55}^{+8.50}\times {10}^{-2}$	1.33
42	${3.27}_{-3.11}^{+6.25}$	${0.67}_{-0.32}^{+0.44}$	${1.06}_{-0.64}^{+8.65}\times {10}^{-2}$	0.97
43	${4.31}_{-2.00}^{+2.90}$	${0.34}_{-0.10}^{+0.13}$	<1.47	2.29
	4.76e	${0.33}_{-0.09}^{+0.04}$	${0.25}_{-0.07}^{+0.72}$
45	${77}_{-17}^{+21}$	${3.87}_{-3.57}^{+4.57}\times {10}^{-2}$	$\lt 3.60\,\times \,{10}^{29}$	1.97
46	${122}_{-34}^{+16}$	<0.06	<416	2.12
47	${2.25}_{-1.41}^{+1.80}$	${0.50}_{-0.15}^{+0.18}$	${2.93}_{-1.58}^{+5.54}\times {10}^{-2}$	2.50
48	${7.80}_{-3.60}^{+6.99}$	${0.29}_{-0.13}^{+0.14}$	<5.13	4.94
	6.50e	${0.33}_{-0.04}^{+0.05}$	${0.19}_{-0.10}^{+0.12}$
49	<11	${0.52}_{-0.18}^{+0.99}$	<0.37	4.22

Notes.

^aFaint sources are excluded (see Section 3.1). ^bWe assume the distance to the SNR to be 6 kpc. ^cThe angular distance from the Fe peak location. ^dWe allow ${{kT}}_{\mathrm{bb}}$ to vary in 0.2–0.5 keV for this source. ^eBest-fit values applied constraints in error ranges on each corresponding SNR region in Table 3.

Download table as:  ASCIITypeset image

Next, we compare the temperature and luminosity of these sources with typical values for CCOs: kT_bb = 0.2–0.5 keV and 10^33-34 erg s⁻¹ (e.g., Pavlov et al. 2004; Gotthelf et al. 2013). We find that four sources, No. 31, 36, 43 and 48 show typical kT_bb and luminosity as CCOs. These sources, however, have relatively large errors of 37%–90% in N_H. Thus, we allow their N_H values to vary in the error ranges of N_H for the corresponding regions in Table 3 for each source. We finally find that all four sources have low luminosities with expected N_H values of the different SNR regions. Therefore we conclude that there is no CCO associated with G344.7–0.1.

Observations of ¹²CO (J = 1–0) line emission at 115.271202 GHz were executed in 2012 May 18th using the NANTEN2 millimeter/submillimeter radio telescope. The telescope is installed at 4865 m altitude of the Atacama Desert in Chile, operated by Nagoya University. We observed an area of 1° × 1° region using the on-the-fly mapping mode with Nyquist sampling. The 4 K cooled Nb superconductor-insulator-superconductor mixer receiver was used for the front end. The system temperature including the atmosphere was ∼200 K in the double-side band. The backend was a digital Fourier-transform spectrometer with 1684 channels or 1 GHz bandwidth, corresponding to a velocity coverage of 2600 km s⁻¹ and a velocity resolution of 0.16 km s⁻¹. After convolving the cube data with a two-dimensional Gaussian kernel of 90'', the final beam size was ∼180'' in FWHM. The absolute intensity was calibrated by observing the standard source IRAS 16293–2422 [(αあるふぁ_J2000, δでるた_J2000) = (16^h32^m233, −24°28'392)] (Ridge et al. 2006). We also observed IRC+10216 every day, and the pointing accuracy was better than 10''. The typical noise fluctuation of the final data set is ∼0.65 K at the velocity resolution of 0.63 km s⁻¹.

The H i line data at 1.4 GHz is from the Southern Galactic Plane Survey (SGPS; McClure-Griffiths et al. 2005). The archival data set was obtained using a combination of the Australia Telescope Compact Array (ATCA) and Parkes radio telescope. The angular resolution is 130'' and the velocity resolution is 0.82 km s⁻¹. The typical noise fluctuation is ∼1.9 K at the velocity resolution of 0.82 km s⁻¹.

To obtain the total proton column density map, we estimate proton column densities in both the molecular and atomic clouds (e.g., Fukui et al. 2012, 2017; Kuriki et al. 2018; Sano et al. 2019). The proton column density of molecular cloud N_p(H₂) can be estimated from the following equations:

$$\begin{eqnarray}&&N({{\rm{H}}}_{2})={X}_{\mathrm{CO}}\cdot W(\mathrm{CO})\ ({\mathrm{cm}}^{-2}),\end{eqnarray} \tag{ B1 }$$

$$\begin{eqnarray}&&{N}_{{\rm{p}}}({{\rm{H}}}_{2})=2\times N({{\rm{H}}}_{2})\ ({\mathrm{cm}}^{-2}),\end{eqnarray} \tag{ B2 }$$

where N(H₂) is the column density of molecular hydrogen in units of cm⁻², X_CO is the CO-to-H₂ conversion factor in units of (K km s⁻¹)⁻¹ cm⁻², and W(CO) is the integrated intensity of CO in units of K km s⁻¹. In the present paper, we used X_CO = 2.0 × 10²⁰ (K km s⁻¹)⁻¹ cm⁻² (Bertsch et al. 1993).

On the other hand, an estimation of the proton column density of atomic cloud N_p(H i) is tricky because we should consider the optical depth of H i. According to Fukui et al. (2015), 91% of atomic hydrogen in local clouds has an optical depth of 0.5 or higher with respect to the H i line emission. Therefore we cannot use the well-known equation assuming the optical depth of H i ≪1 (e.g., Dickey & Lockman 1990):

$$\begin{eqnarray}&&{N}_{{\rm{p}}}({\rm{H}}\,{\rm{I}})=1.823\times {10}^{18}W({\rm{H}}\,{\rm{I}})\ ({\mathrm{cm}}^{-2}),\end{eqnarray} \tag{ B3 }$$

where W(H i) is the integrated intensity of H i in units of K km s⁻¹.

To estimate the optical-depth-corrected proton column density of atomic cloud N_p(H i)', we used a relation between W(H i) and N_p(H i)' presented by Fukui et al. (2017). The authors derived N_p(H i)' as a function of W(H i) using the dust opacity map at 353 GHz obtained from IRAS and Planck data sets (Planck Collaboration et al. 2014) taking into account nonlinear dust properties (e.g., Roy et al. 2013; Okamoto et al. 2017; Hayashi et al. 2019a, 2019b). In the present paper, we derived a conversion factor from W(H i) to N_p(H i)' fitted by the result of Fukui et al. (2017) as a function of W(H i). We obtained averaged value of N_p(H i)' toward G344.7–0.1 is ∼1.2 × 10²¹ cm⁻², which is ∼2.3 times higher than that of N_p(H i) assuming the optically thin case. Finally, we derived the total proton column density map ${N}_{{\rm{p}}}({{\rm{H}}}_{2}+{\rm{H}}\,{\rm{I}})$ combining ${N}_{{\rm{p}}}({\rm{H}}\,{\rm{I}})$' and ${N}_{{\rm{p}}}({{\rm{H}}}_{2})$ whose angular resolution is smoothed to match the FWHM of NANTEN2 beam size of ∼180'' (see Figure 8).

Giacani et al. (2011) argued that G344.7–0.1 is likely associated with an open H i shell at the central velocity of ∼115 km s⁻¹, corresponding to the kinematic distance of 6.3 ± 0.1 kpc. Here, we investigate the cloud association using both the CO and H i data sets.

Figures B1(a) and (c) show the integrated intensity maps of H i and CO at a velocity of ∼115 km s⁻¹, respectively. We found a shell-like structure not only in the H i map, but also in the CO map. The CO clouds nicely surround the SNR, especially for the northern half of the shell. Figures B1(b) and (d) show the position–velocity diagrams of H i and CO, respectively. We found cavity-like structures of H i and CO (dashed lines) in the velocity range from −118.0 km s⁻¹ to −109.1 km s⁻¹, which have a similar diameter to G344.7–0.1 in terms of the decl. range. These cavity-like structures represent an expanding gas motion similar to a "wind-blown shell," originated by strong winds from the progenitor system: e.g., stellar winds from a high-mass progenitor or accretion winds (also known as "disk winds") from a single degenerate progenitor system of Type Ia SNR. For the case of Type Ia SNR G344.7–0.1, the latter origin is favored. The expansion velocity of molecular and atomic clouds ΔでるたV is estimated to be ∼4.5 km s⁻¹, which is consistent with the previous studies of interstellar environments in the Type Ia SNRs (e.g., Tycho, ΔでるたV ∼ 4.5 km s⁻¹, Zhou et al. 2016; N103B, ΔでるたV ∼ 5 km s⁻¹, Sano et al. 2018). In any case, the SNR shock wave travels through the wind-blown shell during the free expansion phase, and is finally decelerated when the shock wave encounters the wind wall. Therefore, the size of the SNR generally coincides with that of the wind-blown shell. Summarizing the above considerations, both the molecular and atomic clouds in the velocity range from −118.0 km s⁻¹ to −109.1 km s⁻¹ are likely associated with the SNR G344.7–0.1, which is consistent with the previous H i study by Giacani et al. (2011).

Zoom In Zoom Out Reset image size