# Discovery of a Wind-Blown Bubble Associated with the Supernova Remnant
G346.6$-$0.2:
A Hint for the Origin of Recombining Plasma

H. Sano Division of Science, National Astronomical Observatory of Japan,
Mitaka, Tokyo 181-8588, Japan<EMAIL_ADDRESS>H. Suzuki Department
of Physics, Konan University, 8-9-1 Okamoto, Higashinada, Kobe, Hyogo
658-8501, Japan  K. K. Nobukawa Faculty of Science and Engineering, Kindai
University, 3-4-1 Kowakae, Higashi-Osaka, 577-8502, Japan  M. D. Filipović
Western Sydney University, Locked Bag 1797, Penrith South DC, NSW 1797,
Australia  Y. Fukui Department of Physics, Nagoya University, Furo-cho,
Chikusa-ku, Nagoya 464-8601, Japan  T. J. Moriya Division of Science, National
Astronomical Observatory of Japan, Mitaka, Tokyo 181-8588, Japan;
<EMAIL_ADDRESS>School of Physics and Astronomy, Faculty of Science,
Monash University, Clayton, Victoria 3800, Australia

###### Abstract

We report on CO and H I studies of the mixed-morphology supernova remnant
(SNR) G346.6$-$0.2. We find a wind-blown bubble along the radio continuum
shell with an expansion velocity of $\sim 10$ km s-1, which was likely formed
by strong stellar winds from the high-mass progenitor of the SNR. The radial
velocities of the CO/H I bubbles at $V_{\mathrm{LSR}}=-82$–$-59$ km s-1 are
also consistent with those of shock-excited 1720 MHz OH masers. The molecular
cloud in the northeastern shell shows a high-kinetic temperature of $\sim 60$
K, suggesting that shock-heating occurred. The H I absorption studies imply
that G346.6$-$0.2 is located on the far side of the Galactic center from us,
and the kinematic distance of the SNR is derived to be $11.1_{-0.3}^{+0.5}$
kpc. We find that the CO line intensity has no specific correlation with the
electron temperature of recombining plasma, implying that the recombining
plasma in G346.6$-$0.2 was likely produced by adiabatic cooling. With our
estimates of the interstellar proton density 280 cm-3 and gamma-ray luminosity
$<5.8\times 10^{34}$ erg s-1, the total energy of accelerated cosmic rays
$W_{\rm p}<9.3\times 10^{47}$ erg is obtained. A comparison of the age–$W_{\rm
p}$ relation to other SNRs suggests that most of the accelerated cosmic rays
in G346.6$-$0.2 have been escaped from the SNR shell.

Supernova remnants (1667); Interstellar medium (847); Cosmic ray sources
(328); Gamma-ray sources (633); X-ray sources (1822)

††software: IDL Astronomy User’s Library (Landsman, 1993), MIRIAD (Sault et
al., 1995), SAS (v19.1.0: Gabriel et al., 2004), ESAS (Kuntz & Snowden, 2008),
HEAsoft (v6.28: Nasa High Energy Astrophysics Science Archive Research Center,
2014), Fermitools (v1.2.23; https://github.com/fermi-lat/Fermitools-
conda/)††facilities: NANTEN2, Mopra, Atacama Pathfinder Experiment (APEX),
Australia Telescope Compact Array (ATCA), Parkes, Molonglo Observatory
Synthesis Telescope (MOST), XMM-Newton, and Fermi-LAT

##  1 Introduction

Over the last three decades, X-ray spectroscopic observations using CCD
detectors allowed us to reveal detailed properties and evolution of shocked
thermal plasma in supernova remnants (SNRs). Because the heavy elements in
both the supernova ejecta and the interstellar medium (ISM) are initially in
low-ionization or neutral before the shock-heating, it is widely accepted that
plasma in young SNRs can be explained as nonequilibrium ionization (NEI,
a.k.a. ionizing plasma). After a long time of evolution, plasma in middle-aged
SNRs reaches collisional ionization equilibrium (CIE). This standard
evolutional scenario, however, has been questioned since the first discovery
of recombining (overionized) plasma where heavy nuclei are stripped of more
electrons than they would be if in the CIE state (Kawasaki et al., 2002).
Although the recombining plasma has been detected from more than a dozen SNRs
(see a review of Yamaguchi, 2020), the origin of the peculiar plasma is still
under debate.

Because the recombining plasma is characterized by lower electron temperature
than the ionization temperature, rapid electron cooling or increasing
ionization state of nuclei is needed to produce the plasma state. Two primary
scenarios are describing the origin of the recombining plasma in SNRs:
adiabatic cooling and thermal conduction scenarios (cf. Yamaguchi, 2020). For
the adiabatic cooling scenario (a.k.a. rarefaction scenario), rapid electron
cooling occurs when the supernova shockwaves breakout from a dense
circumstellar matter (CSM) into the surrounding low-density region such as a
wind-blown bubble (e.g., Itoh & Masai, 1989; Masai, 1994). For the thermal
conduction scenario, on the other hand, rapid electron cooling is caused by
shock-interactions with cold/dense molecular clouds via thermal conduction
(e.g., Kawasaki et al., 2002, 2005). In fact, the lower electron temperature
of recombining plasma is found toward the shocked molecular clouds in the SNRs
IC443, W28, W44, and W49B (Matsumura et al., 2017; Okon et al., 2018, 2020;
Sano et al., 2021). To distinguish the two primary scenarios, therefore, it is
essential to investigate the spatial distribution of plasma conditions and
their relation to the ISM.

G346.6$-$0.2 is one of the mixed-morphology SNRs that are classified as radio
shell morphology with centrally filled thermal X-rays (Rho & Petre, 1998). The
small apparent diameter of radio shell $\sim 8^{\prime}$ is consistent with a
large distance from us ($\sim 5.5$–11 kpc, Koralesky et al., 1998; Dubner et
al., 1993) and its old age ($\sim 12$–16 kyr, Sezer et al., 2011; Yamauchi et
al., 2013; Auchettl et al., 2017). The thermal X-rays are well reproduced by
the recombining plasma model (Yamauchi et al., 2013; Auchettl et al., 2017).
According to Auchettl et al. (2017), the recombining plasma likely arises from
adiabatic cooling by comparing time scales of thermal conduction and adiabatic
cooling. The authors also found additional hard X-rays comprising either a
power-law component with a photon index of $\sim 2$ or a thermal component
with a temperature of $\sim 2.0$ keV. In either case, spatially resolved X-ray
spectroscopy presented no significant differences in the absorbing column
densities, electron temperature of recombining plasma, and the photon indexes
toward the SNR (Auchettl et al., 2017).

G346.6$-$0.2 is also believed to be interacting with dense molecular clouds.
Koralesky et al. (1998) discovered several shock-excited 1720 MHz OH masers
along with the southern shell of the SNR at a velocity range from $-79.3$ to
$-74.0$ km s-1, implying that the SNR is physically associated with molecular
clouds at the same velocities. Subsequent infrared studies revealed the
presence of shocked H2 emission in the southern rim and possibly in the
northern shell (Reach et al., 2006; Andersen et al., 2011). Nevertheless,
there are no radio-line studies using CO and H I that can trace the bulk
masses of molecular and atomic clouds. Moreover, hadronic gamma-rays—which are
produced by interactions between cosmic-ray protons and ISM protons—have never
been detected from the SNR (Ergin & Ercan, 2012).

In the present paper, we report on detailed CO and H I studies using the radio
telescopes and interferometers of NANTEN2, Mopra, Atacama Pathfinder
Experiment (APEX), Australia Telescope Compact Array (ATCA), and Parkes to
reveal the origin of the recombining plasma from G346.6$-$0.2. Section 2
describes datasets of CO, H I, radio continuum, X-rays, and gamma-rays.
Section 3 comprises five subsections: Section 3.1 presents distributions of
X-ray and radio continuum. Sections 3.2–3.4 show CO and H I distributions and
their physical properties. Section 3.5 provides results on the gamma-ray
analysis. Discussion and conclusions are given in Sections 4 and 5,
respectively.

Table 1: Summary of XMM-Newton archive data in G346.6$-$0.2

|  |  |  |  | Exposure
---|---|---|---|---|---
Obs. ID | $\alpha_{\mathrm{J2000}}$ | $\delta_{\mathrm{J2000}}$ | Start Date | End Date | MOS1 | MOS2 | pn
| (h m s) | ($\arcdeg$ $\arcmin$ $\arcsec$) | (yyyy-mm-dd hh:mm:ss) | (yyyy-mm-dd hh:mm:ss) | (ks) | (ks) | (ks)
0654140101 | 17 10 17.00 |  $-40$ 10 22.1 | 2011-03-15 11:51:02 | 2011-03-15 20:13:02 | 29.1 | 28.4 | 27.2
0782080101 | 17 10 09.70 |  $-40$ 11 32.5 | 2016-09-04 16:07:23 | 2016-09-05 23:14:02 | 97.5 | 95.2 | 92.6

Note. — All exposure times represent the flare-filtered exposure.

##  2 Datasets

###  2.1 CO, H I, and Radio Continuum

Observations of 12CO($J$ = 2–1) line emission at 230 GHz were conducted from
August to November in 2008 using the NANTEN2 4-m millimeter/sub-millimeter
radio telescope belonging to Nagoya University, which has been installed at
Pampa La Bola (4865-m above sea level) in northern Chile. We carried out a
Nyquist sampled on-the-ﬂy mapping which covered an area of 2.75 degree2
containing both the SNRs RX J1713.7$-$3946 and G346.6$-$0.2. Most of the
datasets have been published in several papers for RX J1713.7$-$3946 (Fukui,
2008; Fukui et al., 2012; Sano et al., 2010, 2013; Maxted et al., 2012, 2013).
The frontend was a 4 K cooled Nb superconductor-insulator-superconductor mixer
receiver. The typical system temperature including atmosphere was $\sim 250$ K
in the single sideband. The backend was an acoustic optical spectrometer with
2048 channels, providing a velocity coverage of 390 km s-1 with a velocity
resolution of 0.38 km s-1 at 230 GHz. After convolution using a two-
dimensional Gaussian function of $45^{\prime\prime}$, the final beam size was
$\sim 90^{\prime\prime}$ in the Full-Width-Half-Maximum (FWHM). The absolute
intensity was calibrated by observing Orion-KL [($\alpha_{\mathrm{J2000}}$,
$\delta_{\mathrm{J2000}}$) $=$ ($05^{\mathrm{h}}35^{\mathrm{m}}14\fs 52$,
$-05\arcdeg 22\arcmin 28\farcs 2$)] (Schneider et al., 1998). The pointing
accuracy was achieved to be better than $\sim 15^{\prime\prime}$ through two-
hourly observations of Jupiter. The typical noise fluctuations toward
G346.6$-$0.2 are $\sim 0.12$ K at a velocity resolution of 1 km s-1.

We also used archival datasets of 12CO($J$ = 1–0) and 13CO($J$ = 1–0, 2–1)
line emission to derive the physical properties of molecular clouds: e.g.,
mass, kinetic temperature, and number density of molecular hydrogen. The
12CO($J$ = 1–0) and 13CO($J$ = 1–0) data are from the Mopra Southern Galactic
Plane CO Survey Data Release 3 (DR3, Braiding et al., 2018) using the Mopra
22-m radio telescope, and the 13CO($J$ = 2–1) data are from the Structure,
Excitation and Dynamics of the Inner Galactic Interstellar Medium survey
(SEDIGISM; Schuller et al., 2021) using the APEX 12-m radio telescope. The
angular resolution is $\sim 36^{\prime\prime}$ for the 12CO($J$ = 1–0) and
13CO($J$ = 1–0) data; and $\sim 30^{\prime\prime}$ for the 13CO($J$ = 2–1)
data. To compare the 12CO($J$ = 2–1) data obtained with NANTEN2, we smoothed
all the datasets to match the FWHM of $90^{\prime\prime}$ using a two-
dimensional Gaussian function. The typical noise fluctuations are $\sim 0.18$
K for the 12CO($J$ = 1–0) and 13CO($J$ = 1–0) data; and $\sim 0.06$ K for the
13CO($J$ = 2–1) data at a velocity resolution of 1 km s-1 for each.

The H I line data at 1.4 GHz were provided by the Southern Galactic Plane
Survey (SGPS; McClure-Griffiths et al., 2005) using ATCA combined with the
Parkes 64-m radio telescope. The combined beam size is
$130^{\prime\prime}\times 130^{\prime\prime}$ and a velocity resolution of
0.82 km s-1. The typical noise fluctuations are $\sim 1.3$ K at a velocity
resolution of 1 km s-1.

The radio continuum data at 843 MHz are from the Molonglo Observatory
Synthesis Telescope (MOST) supernova remnant catalogue (MSC, Whiteoak & Green,
1996). The angular resolution is $\sim 43^{\prime\prime}$ and the typical
noise fluctuations are $\sim 2$ mJy beam-1.

###  2.2 X-rays

We used archival datasets obtained with XMM-Newton to present a map of X-ray
recombining plasma in G346.6$-$0.2, for which the observation IDs are
0654140101 (PI: C.-Y. Ng) and 0782080101 (PI: K. Auchettl). We used the XMM-
Newton Science Analysis System (SAS, Gabriel et al., 2004) version 19.1.0 and
HEAsoft version 6.28 to analyze both the EPIC-MOS and EPIC-pn datasets which
were obtained using the full-frame mode with the thick filter. Table 1 lists
the details of the observations. We reprocessed the Observation Data Files
(ODF) for each pointing, following a standard procedure provided as the XMM-
Newton Extended Source Analysis Software (ESAS, Kuntz & Snowden, 2008). The
effective exposures after filtering soft proton flares are also shown in Table
1. The total exposure time is $\sim 370$ ks. To subtract the Quiescent
Particle Background (QPB) from the X-ray map, we run “mos-$/$pn-back” and
“mos-$/$pn-filter” procedures. We also used the procedure “merge_comp_xmm” to
combine the two-pointing data. Note that we excluded some CCD chips or
quadrants which are affected by strong stray light (see also Figure 2 of
Auchettl et al., 2017). We then applied an adaptive smoothing by using the
“adopt_merge” procedure, where the pixel size and smoothing counts were
$6^{\prime\prime}$ and 300 counts, respectively. Finally, we produced a QPB-
subtracted, exposure-corrected, and adaptively smoothed image in the energy
band of 0.5–7.0 keV.

###  2.3 Gamma-Rays

In order to get information of particle acceleration, we also studied the
gamma-ray emission around this source with Fermi-LAT. We extracted all the
available data for the circle region with a radius of $20^{\circ}$ centered at
the source position from the Pass 8
database111https://fermi.gsfc.nasa.gov/cgi-bin/ssc/LAT/LATDataQuery.cgi. The
energy range for the data extraction is 100 MeV to 300 GeV.

The tool and databases used in the data reduction/analysis are Fermitools
(v1.2.23)222https://github.com/fermi-lat/Fermitools-conda/, the Instrumental
Response File version P8R3_SOURCE_V2, the Fermi source list gll_psc_v22.fit
(4FGL catalog), the Galactic diffuse background model gll_iem_v07.fits, and
the isotropic background model (instrumental and extragalactic)
iso_P8R3_SOUCE_V2.txt.

The standard event selection is applied to the retrieved data set: from the
SOURCE class, both FRONT and BACK section events are extracted (evclass=128
evtype=3). Events with zenith angles larger than 90∘ are rejected from the
analysis in order to prevent contamination from the Earth’s bright limb. The
energy dispersion correction is enabled for all the model components but the
isotropic background
model333https://fermi.gsfc.nasa.gov/ssc/data/analysis/documentation/Pass8_edisp_usage.html.

Figure 1: Left panel: Map of the XMM-Newton X-ray flux toward the SNR
G346.6$-$0.2 ($E$: 0.5–7.0 keV). The superposed contours indicate the 843 MHz
radio continuum from the MOST Supernova Remnant Catalogue (Whiteoak & Green,
1996). The lowest contour level and the contour intervals are 20 and 40 mJy
beam-1, respectively. The crosses represent the positions of 1720 MHz OH
masers (Koralesky et al., 1998). Right panel: Radial profiles of the radio
continuum, centered at ($l$, $b$) = ($346\fdg 62$, $-0\fdg 22$). The small
dots represent all the data points for the radio continuum. The black steps
with large filled circles indicate averaged values of radio continuum at each
annulus. The green line represents the three-dimensional Gaussian distribution
and the orange line represents its projected distribution derived by the
least-squares fitting (see the text). The vertical dashed lines indicate best-
fit values of radius and ranges of shell thickness.

##  3 Results

###  3.1 Overview of X-rays and Radio Continuum

Figure 1 left panel shows the overlay map of X-rays and radio continuum toward
G346.6$-$0.2. As presented in previous studies, the diffuse X-ray emission
produced by recombining plasma shows the centrally filled distribution inside
the radio continuum shell (Yamauchi et al., 2013; Auchettl et al., 2017). The
highest value of X-ray flux found in the center and northwestern regions,
while the northeast, west, and southeast regions inside the shell are dim in
the X-rays. Note that this trend is not caused by interstellar absorption
because the spatially resolved X-ray spectroscopy indicated almost uniform
absorbing column densities within the shell (Auchettl et al., 2017). We also
find that the positions of the shock-excited OH masers spatially correspond to
the outer boundary of the diffuse X-ray emission. There are a dozen of X-ray
point sources, almost all of which have been analyzed by Auchettl et al.
(2017) and will not be discussed further in this paper.

To derive the apparent diameter and shell thickness of the radio shell
quantitatively, we fitted its radial profile using a three-dimensional
spherical shell with a Gaussian function $F(r)$:

$\displaystyle F(r)=A\exp[-(r-r_{0}\bigr{)}^{2}/2\sigma^{2}]$ (1)

where $A$ is a normalization of the Gaussian function, $r_{0}$ is the shell
radius in units of arcmin, and $\sigma$ is the standard deviation of the
Gaussian function in units of arcmin. We first fitted the radial profile by
changing an original position around the geometric center of the SNR. This
process gave the central position of the shell to be ($l$, $b$) =
(346$\fdg$62, $-$0$\fdg$22) with the minimum chi-square value of the least-
squares fitting. Figure 1 right panel shows the radial profiles of the radio
continuum and X-rays. We derived the shell radius $r_{0}$ of $3\farcm 29\pm
0\farcm 05$ and thickness of $1\farcm 41\pm 0\farcm 11$ as the best-fit
parameters, where the shell thickness is defined as the FWHM of the Gaussian
function or $2\sigma\sqrt{2\ln 2}$.

Figure 2: Velocity channel distributions of the NANTEN2 12CO($J$ = 2–1) (top
panels) and ATCA & Parkes H I (bottom panels). The superposed contours and
crosses are the same as shown in Figure 1. Each panel shows CO or H I
distributions every 6 km s-1 in a velocity range from $-88$ to $-58$ km s-1.
The CO clouds in the northeast, south, and possibly west are also enclosed by
dashed yellow circles. Figure 3: Integrated intensity maps and
position–velocity (p–v) diagrams of 12CO($J$ = 2–1) (top panels) and H I
(bottom panels). The integration range is from $-82$ to $-59$ km s-1 in the
velocity for each intensity map; and from 346$\fdg$54 to 346$\fdg$70 in the
galactic longitude for each p–v diagram. The solid contours represent the
radio continuum whose contour levels are the same as shown in Figure 1. The
lowest contour level and the contour intervals are 12 and 3 K km-1 for the CO
intensity map; 800 and 80 K km-1 for the H I intensity map; 0.08 and 0.04 K
degree for the CO p–v diagram; 5.5 and 0.5 K degree for the H I p–v diagram.
The red crosses represent the positions of 1720 MHz OH masers (Koralesky et
al., 1998). The CO peaks “NE-shell” and “Reference” discussed in Section 3.3
are also indicated in (a). The dashed circles in the p–v diagrams indicate the
boundaries of the CO and H I cavities (see the text).

###  3.2 Distributions of CO and H I

Figure 2 shows the velocity channel maps of CO and H I toward G346.6$-$0.2. In
the present paper, we focus on only the velocity range from $-88$ to $-58$ km
s-1, which is roughly consistent with the radial velocities of the shock-
excited OH masers (Koralesky et al., 1998). We find several molecular clouds
overlapping with the radio shell, especially toward the northeast, south, and
possibly in the west as shown by dashed yellow circles. The south CO clouds
show a good spatial correspondence with the positions of the shock-excited OH
masers. In the H I channel maps, we find a cavity-like distribution toward the
radio shell in the velocity range from $-82$ to $-58$ km s-1. The rich H I
clouds are located along with the northwestern half of the SNR, while the
southeastern half shows a blowout structure.

Figures 3a and 3c show the integrated intensity maps of CO and H I at
$V_{\mathrm{LSR}}=-82$–$-58$ km s-1. The spatial relation between the radio
shell and CO/H I clouds is clearer than the velocity channel maps. We note
that the CO intensity of the south clouds overlapping with the shock-excited
OH masers are weaker than that of the molecular cloud laying on the edge of
the northeastern shell (refer to as “NE-shell cloud”). Moreover, there are no
dense CO/H I clouds toward the center of the SNR where the X-ray recombining
plasma is strongly detected.

Figures 3b and 3d show the position–velocity (p–v) diagrams of CO and H I. We
find a cavity-like structure in the p–v diagram of H I, whose velocity range
is from $-82$ to $-58$ km s-1. Note that all the shock-excited OH masers are
located inside the velocity range of the cavity. It is noteworthy that the
spatial extent of the cavity in the Galactic longitude direction is roughly
consistent with the apparent diameter of the radio shell. Moreover, the shock-
excited OH masers are located on the low-density regions (or surface) of the
molecular clouds in the velocity space (see Figure 3b).

###  3.3 LVG Analysis

To estimate the kinetic temperature and density of molecular clouds, we
performed the Large Velocity Gradient analysis (LVG, e.g., Goldreich & Kwan,
1974; Scoville & Solomon, 1974). The LVG analysis calculates the radiative
transfer of multi-transitions of molecular line emission, assuming a
spherically isotropic cloud with uniform kinetic temperature, photon escape
probability, and a velocity gradient of $dv/dr$. Here, $dv$ is the Half-Width-
Half-Maximum (HWHM) of CO line profiles obtained by the least-square fitting
using a single Gaussian function, and $dr$ is a cloud radius. To test the
shock-heating and/or compression, we selected two molecular clouds which are
significantly detected both in the 12CO($J$ = 2–1) and 13CO($J$ = 1–0, 2–1)
line emission as well as free from extra heating sources such as infrared
and/or young stellar objects444Although the southern radio shell is
overlapping with both the OH masers and weak CO emission, there are no
molecular clouds that are detected in all molecular lines as well as free from
any extra heating sources.. One of the clouds is the NE-shell cloud which is a
candidate for the post-shocked cloud. The other was selected as a reference
which is located outside of the radio shell (refer to as “reference cloud”).
Therefore, it is expected that the shock-heating and/or compression will be
seen only in the NE-shell cloud if the cloud is interacting with the SNR.

Figure 4: Top panels: CO intensity profiles toward the CO peaks of the NE-
shell and Reference clouds. Each CO data was smoothed to match the beam size
$90^{\prime\prime}$ and the velocity resolution of 0.5 km s-1. Bottom panels:
Results of LVG analysis on the number density of molecular hydrogen $n$(H2)
and the kinetic temperature $T_{\mathrm{kin}}$ for each CO peak. The red lines
and blue dash-dotted lines indicate the intensity ratios of 13CO($J$ = 2–1) /
13CO($J$ = 2–1) and 12CO($J$ = 2–1) / 13CO($J$ = 1–0), respectively. The
shaded areas in red and blue indicate $1\sigma$ error regions for each
intensity ratio. Yellow crosses represent the best-fit values of $n$(H2) and
$T_{\mathrm{kin}}$ for each CO peak. Figure 5: (a–c) H I and CO spectra of
three rectangular regions 1–3 toward G346.6$-$0.2. Top half of each panel
shows H I profiles of source region (black), background region (red), and the
residual (blue). Bottom half of each panel indicates H I absorption profiles
of $e^{-\tau}$ (blue) and 12CO($J$ = 1–0) line emission (green). The gray
shaded areas represent velocity ranges that show H I absorption features in
all three regions. (d) Integrated intensity of H I superposed on the radio
continuum contours. The integrated velocity range is from $-95$ to $-85$ km
s-1. The contour levels of radio continuum are the same as shown in Figure 1.
The three rectangular regions 1–3 are also indicated.

Figure 4 upper panels show the CO spectra toward the two clouds. We note that
the 12CO($J$ = 2–1) profile of the NE-shell cloud shows multiple peaks at
$\sim-86$, $-80$, and $-70$ km s-1 with a broad component at the velocity
range approximately from $-90$ to $-60$ km s-1. We estimated $dv/dr=0.5$ km
s-1 pc-1 for the NE-shell cloud at the $\sim 70$ km s-1 component, and 1 km
s-1 pc-1 for the reference cloud, assuming the largest expected distance of 11
kpc. In the case of the near-side distance at 5.5 kpc, the values of $dv/dr$
will decrease the only factor of two and will not affect the results
significantly. We also used the abundance ratio of [12CO/H2] = $5\times
10^{-5}$ (Blake et al., 1987) and the abundance ratio of isotopes
$[^{12}\mathrm{CO}/^{13}\mathrm{CO}]=33$ (Langer & Penzias, 1990).

Figure 4 bottom panels show the results of the LVG analysis toward the NE-
shell and reference clouds, which describe a relation between the kinetic
temperature $T_{\mathrm{kin}}$ and the number density of molecular hydrogen
$n(\mathrm{H}_{2})$. We derived the best-fit parameters of
$T_{\mathrm{kin}}=64^{+18}_{-14}$ K and $n(\mathrm{H}_{2})=790^{+100}_{-70}$
cm-3 for the NE-shell cloud, and $T_{\mathrm{kin}}=19^{+3}_{-3}$ K and
$n(\mathrm{H}_{2})=1000^{+70}_{-50}$ cm-3 for the reference cloud. The values
of $n(\mathrm{H}_{2})$ are roughly consistent with each other ($\sim 1000$
cm-3), while the kinetic temperature of the NE-shell cloud is more than three
times as high as that of the reference cloud.

###  3.4 H I Absorption Studies

To investigate the distance of G346.6$-$0.2, we carried out H I absorption
studies. Following the latest method described by Leahy & Tian (2010) and
Ranasinghe & Leahy (2017), we selected three regions 1, 2, and 3 toward the
SNR (Figure 5d). Figures 5a–5c show the averaged H I source, background, and
residual (or $e^{-\tau}$) profiles as well as the CO spectra toward the
regions 1, 2, and 3, respectively. We find significant absorption features of
all H I spectra at three velocity ranges of $V_{\mathrm{LSR}}=-95$–$-85$,
$-50$–$-30$, and $-20$–$-2$ km s-1. The three velocity ranges clearly show
negative residuals (or $e^{-\tau}<1$) without strong CO line emission,
suggesting that the absorption profiles are not caused by the self-absorption
effect due to cold/dense clouds. Moreover, the H I integrated intensity maps
of three velocity ranges show intensity dips of H I toward the radio shell.
Figure 5d shows the H I integrated intensity map at the velocity range from
$-95$ to $-85$ km s-1. The H I intensity dips show a nice spatial
correspondence with the radio bright shell. This means that the H I absorption
features due to the radio continuum emission from the SNR found at least the
velocity range from $-95$ to $-2$ km s-1. The interpretation of this
observational result will be discussed in Section 4.2 in detail.

###  3.5 Upper Limit of GeV Gamma-Ray Flux

Here, gamma-ray emission at GeV energies from the source was examined. The
analysis region was a 20∘-radius circle centered at the SNR position. In order
to get the energy flux of the source, a binned maximum likelihood analysis on
the spatial and energy distributions of the data was conducted. The spatial
bin size was $0^{\circ}_{\cdot}2\times 0^{\circ}_{\cdot}2$. The data were
divided into 35 logarithmic energy bins in the energy range of 100 MeV–300
GeV. All the sources in the 4FGL catalog within our analysis region were
considered in the likelihood analysis, and the spectral parameters of only
those within a $8^{\circ}_{\cdot}5$-radius circle were treated as free
parameters. For the sources included in the 4FGL catalog, the spectral models
in the catalog were used. For G346.6$-$0.2 itself, a simple power-law model
was assumed.

The maximum likelihood analysis was conducted to search for the model
configuration which best represents the data. This procedure was repeated with
decreasing fit tolerance values until the difference between the resultant
likelihood and that in the last trial becomes less than unity. Then, we
obtained the best-fit source flux. The test-statistic (TS) for this source
was, however, found to be 12.4, which corresponds to a $\approx 3.5\sigma$
detection. Thus, we concluded that the source detection was marginal or
insignificant, and derived only the upper limit of the energy flux as
described below.

In the calculation of the upper limit of the energy flux, we considered
systematic uncertainties associated with following components: Fermi-LAT’s
effective area ($\pm 3\%$) and point spread function ($\pm
5\%$)555https://fermi.gsfc.nasa.gov/ssc/data/analysis/LAT_caveats.html, and
the Galactic diffuse background model ($\pm 6\%$; e.g., Abdo et al. 2009;
Castro & Slane 2010; Tanaka et al. 2011). We ran the maximum likelihood
analysis several times considering these potential uncertainties, and obtained
the energy fluxes and errors of the source in individual runs. The largest
value among these was used for estimation of the upper limit. The resultant
95% upper limit of the energy flux is $4.0\times 10^{-12}$ erg s-1 cm-2 in the
1–100 GeV energy range, which is converted to $1.4\times 10^{34}$ erg s-1 at
5.5 kpc or $5.8\times 10^{34}$ erg s-1 at 11 kpc.

##  4 Discussion

###  4.1 Molecular and Atomic Clouds Associated with the SNR G346.6$-$0.2

The previous ISM studies of G346.6$-$0.2 suggested the presence of shocked
molecular clouds by observing the shock-excited 1720 MHz OH masers at
$\sim-70$ km s-1 and the near-infrared shocked H2 emission (Koralesky et al.,
1998; Reach et al., 2006; Andersen et al., 2011). However, the
shock–interacting neutral gaseous medium—molecular and atomic hydrogen
clouds—has not been reported due to a lack of comprehensive studies using CO
and H I emission lines. In the present section, we argue that both the
molecular and atomic clouds at $V_{\mathrm{LSR}}=-82$–$-59$ km s-1 are
physically associated with the SNR G346.6$-$0.2.

First, we argue that the cavity-like structure in the p–v diagram of H I
(partially CO as well) provides strong support for the shock–interactions with
the molecular and atomic clouds at $V_{\mathrm{LSR}}=-82$–$-59$ km s-1.
Because such cavity-like structure toward an SNR corresponds to an expanding
gas motion, also called the “wind-blown bubble”, which is thought to be formed
by strong winds from the progenitor system of the SNR: e.g., stellar wind from
the high-mass progenitor or accretion wind from the progenitor system
containing a white dwarf and a companion star (e.g., Koo et al., 1990; Koo &
Heiles, 1991; Hachisu et al., 1999a, b). It is noteworthy that the size of
such a wind-blown bubble is expected to be consistent with that of the SNR
shell because the shock propagation time inside the bubble is very short due
to a much-lower density (e.g., Weaver et al., 1977). Moreover, the expansion
velocities of wind-blown bubbles were observationally derived to be $\sim
3$–13 km s-1 in the several Galactic and Magellanic SNRs (e.g., Landecker et
al., 1989; Fukui et al., 2012; Zhou et al., 2016; Sano et al., 2017, 2018,
2019; Kuriki et al., 2018). In the case of G346.6$-$0.2, the wind-blown bubble
shows the expanding velocity of $\sim 10$ km s-1 and the size of which is
roughly consistent with that of the radio continuum shell (see Figure 3).
Although the progenitor type of G346.6$-$0.2 is still under debate, the core-
collapse origin is thought to be favored (Yamauchi et al., 2013; Auchettl et
al., 2017). We, therefore, propose that the wind-blown bubble at
$V_{\mathrm{LSR}}=-82$–$-58$ km s-1 was formed by the strong stellar wind from
the high-mass progenitor and is now interacting with supernova shockwaves.

The good correspondences between the CO clouds and shock-excited 1720 MHz OH
masers are also consistent with this scenario. In fact, the positions of OH
masers spatially coincide with the CO clouds in the southern shell (see
Figures 2, 3a, and 3c). The OH masers also lie on the edge of the CO clouds in
the velocity space, whose center velocities are in the range of the wind-blown
bubble (see Figures 3b and 3d). Because moderate densities of
$10^{3}$–$10^{5}$ cm-3 and temperature of 50–125 K are needed to efficiently
emit the shock-excited 1720 MHz OH maser (Elitzur, 1976), the slight velocity
offsets between the CO intensity peaks (the densest part of the clouds) and
the OH masers are naturally expected (see Figure 3b).

The high-kinetic temperature in the NE-shell cloud provides alternative
evidence for the shock-cloud interaction. In Figure 4, we can find that the
kinetic temperature of the NE-shell cloud (= post-shocked cloud) is more than
three times as high as that of the reference cloud (= pre-shocked cloud).
Because there are no extra heating sources except for the supernova shocks,
the high temperature of the NE-shell can be understood as the shock-heating.
The kinetic temperature of $\sim 60$ K in the NE-shell cloud is also
consistent with that of shock-heated molecular clouds in the similar mixed-
morphology SNRs W44, IC443, and W28 ($T_{\mathrm{kin}}\sim 40$–80 K, Seta et
al., 1998; Arikawa et al., 1999; Anderl et al., 2014). In addition, the broad
velocity component of the NE-shell cloud at $V_{\mathrm{LSR}}\sim-90$–$-60$ km
s-1 possibly corresponds to the line-broadening due to the shock acceleration
(e.g., Wootten, 1977, 1981; Denoyer, 1979). Further high spatial resolution
and high sensitivity CO observations using the Atacama Large
Millimeter/submillimeter Array (ALMA) will confirm the line-broadening in
detail. In light of these considerations, we conclude that the CO/H I clouds
as well as the wind-blown bubble in the velocity range from $-82$ to $-59$ km
s-1 are physically associated with the SNR G346.6$-$0.2.

###  4.2 Distance and Age

Based on the physical association between the shockwaves and the stellar wind
bubble, we here discuss the distance and age of G346.6$-$0.2. We first derived
the systemic velocity of the stellar wind bubble as
$V_{\mathrm{LSR}}=-72^{+13}_{-10}$ km s-1 from the p–v diagram of H I (Figure
3d). Here, the systemic velocity is defined as the velocity that shows the
largest size of the cavity, and its errors represent to the minimum and
maximum velocities of the wind bubble. By adopting the Galactic rotation curve
model with the IAU-recommended values of $R_{0}=8.5$ kpc and $\Theta_{0}=220$
km s-1 (Kerr & Lynden-Bell, 1986; Brand & Blitz, 1993), we obtained the
kinematic distance of G346.6$-$0.2 as $5.4^{+0.3}_{-0.5}$ kpc for the near
side distance and $11.1^{+0.5}_{-0.3}$ kpc for the far side distance.

Next, we argue that the far side distance would be appropriate for the
distance to G346.6$-$0.2. In case the SNR is located in the near side
distance, we can find the H I absorption feature only in the velocity range
from $-72$ to 0 km s-1 because the absorption line can be seen in the
foreground gas with respect to the SNR. However, our results show that the
absorption features due to the SNR are seen in the velocity range from $-95$
to $-2$ km s-1 (see Section 3.4 and Figure 5). We, therefore, conclude that
the SNR G346.6$-$0.2 is located on the far side of the Galactic center from us
and the appropriate distance is to be $11.1^{+0.5}_{-0.3}$ kpc666Although the
distance is slightly inconsistent with the previous distance of $\sim 9$ kpc
derived using the $\Sigma$–$D$ relation (Dubner et al., 1993), it does not
matter considering a large scatter of the relation (e.g., Pavlovic et al.,
2014).. We then derived the radius of the SNR to be $10.6^{+0.5}_{-0.3}$ pc.

To estimate the age of G346.6$-$0.2, we calculated the electron density
$n_{\mathrm{e}}$ within the stellar wind bubble. According to Auchettl et al.
(2017), the electron density with a plasma filling factor $f$ was derived to
as $1.4f^{-0.5}(d/8.3\>\mathrm{kpc})^{-0.5}$ cm-3 using the emission measure
of the recombining plasma component as well as assuming the electron to ion
ratio and the plasma emitting volume (see Auchettl et al. 2017 for details).
We then obtained the electron density $n_{\mathrm{e}}\sim 1.2$ cm-3 at the SNR
distance of 11.1 kpc. Because the recombination timescale
$n_{\mathrm{e}}t=(5.3\pm 0.6)\times 10^{11}$ cm-3 s, the elapsed time $t$
after producing the recombining plasma can be estimated to $14\pm 2$
kyr777Note that the plasma age of the SNR is expected to have a systematic
error within a factor of four in addition to its statistical error of 2 kyr
(cf. Suzuki et al., 2021).. This values is roughly consistent with the
previous age estimation of the SNR (Sezer et al., 2011; Yamauchi et al., 2013;
Auchettl et al., 2017). In the present paper, we use $\sim$14 kyr as the age
of the SNR G346.6$-$0.2888According to Itoh & Masai (1989), adiabatic cooling
occurs a few hundred years after a supernova explosion. We thus assumed that
the plasma age is roughly consistent with the SNR age..

Figure 6: (a) Same 12CO($J$ = 2–1) intensity map as shown in Figure 3(a), but
the superposed white contours indicate the XMM-Newton X-ray flux. The lowest
contour level and the contour intervals of X-rays are 12 and 3 counts s-1
degree-2, respectively. The shell boundary of radio continuum is also shown in
the yellow contour. The red crosses indicate the positions of 1720 MHz OH
masers (Koralesky et al., 1998). (b) Same as (a) except that the colored image
is replaced by the XMM-Newton X-ray flux as shown in Figure 1. The green boxes
1–6 represent the extracted areas of X-ray spectra as defined by Auchettl et
al. (2017). (c) Scatter plot between the electron temperature
$kT_{\mathrm{e}}$ and the peak integrated intensity of 12CO($J$ = 2–1) line
emission. The vertical error bars represent the standard deviations of CO
intensities for each region.

###  4.3 Origin of the Recombining Plasma

Auchettl et al. (2017) proposed that the recombining plasma in G346.6$-$0.2
was likely formed by adiabatic cooling because the adiabatic cooling timescale
of $\sim 12$ kyr is compatible with the age of the SNR. The authors also
derived the cooling timescale for thermal conduction to be $\sim 500$ kyr, and
hence concluded that thermal conduction is most likely not responsible for
producing the recombining plasma. In this section, we argue that our ISM
results also support the adiabatic cooling scenario as the origin of the
recombining plasma in G346.6$-$0.2.

Figures 6a and 6b show the overlay maps of the CO intensity and the X-ray
flux. We find a clear spatial anti-correlation between the recombining plasma
and shock-interacting molecular clouds especially toward the center of the
SNR. Figure 6c shows the scatter plot between the electron temperature
$kT_{\mathrm{e}}$ and the peak integrated intensity of 12CO($J$ = 2–1) line
emission for the regions 1–6 which were defined by Auchettl et al. (2017).
Although a negative correlation between the CO intensity and $kT_{\mathrm{e}}$
will be expected for the thermal conduction scenario (Matsumura et al., 2017;
Okon et al., 2018, 2020; Sano et al., 2021), we could not find such a
correlation (correlation coefficient $\sim-0.09$). It is noteworthy that the
regions containing the shocked molecular clouds do not always exhibit the
lower $kT_{\mathrm{e}}$ values. These results, therefore, disfavor the pure
thermal conduction scenario as the formation mechanism of the recombining
plasma in G346.6$-$0.2.

We also would like to emphasize that the supernova explosion inside the
stellar wind-bubble can naturally explain the adiabatic cooling scenario as
the origin of recombining plasma in G346.6$-$0.2. Before the supernova
explosion, a high-mass progenitor loses its gaseous envelope via stellar
winds. The stellar winds (or mass loss events) can form not only a wind-blown
bubble but also a dense CSM in the vicinity of the progenitor (e.g., Georgy et
al. 2013 and references therein), and then the supernova explosion occurs
inside the dense CSM. According to Moriya (2012), the dense CSM around red
supergiants and Type IIn supernova progenitors can establish the CIE state of
the plasma soon after the supernova explosion, and its remnant can evolve to
the recombining plasma SNR via adiabatic cooling when the supernova shock
breaks out of the CSM. In the case of G346.6$-$0.2, the presence of the
stellar wind bubble is consistent with the scenario. If we take the
relationship between the stellar mass and the wind-blown bubble size estimated
by Chen et al. (2013), the progenitor mass of G346.6$-$0.2 is estimated to be
$16~{}\mathrm{M_{\odot}}$ for the wind-blown bubble size of 10.6 pc. The
progenitor mass is consistent with a massive red supergiant progenitor
(Smartt, 2015) which is suggested to be a progenitor of recombining plasma
SNRs (Moriya, 2012). Further sensitive observations with high-spatial
resolutions using ALMA will possibly unveil the remnant of the dense CSM
toward the center of the SNR.

Figure 7: Correlation plot between the age of SNRs and the total energy of
cosmic-rays $W_{\mathrm{p}}$ (Sano et al., 2021). The green line indicates the
linear regression of the double-logarithmic plot applying the least-squares
fitting for data points with the ages of SNRs below 6 kyr.

###  4.4 Total Energy of Accelerated Cosmic-Rays

SNRs are believed to be the primary sources of Galactic cosmic rays, mainly
consisting of protons, up to at least the energy of $\sim 3$ PeV via the
diffusive shock acceleration (DSA, e.g., Bell, 1978; Blandford & Ostriker,
1978). Although the conventional value of total energy of accelerated cosmic
rays $W_{\mathrm{p}}$ is thought to be $\sim 10^{49}$–$10^{50}$ erg per a
single supernova explosion, the observational constraint for $W_{\mathrm{p}}$
was insufficient because of lack of unified identification and/or
quantification for shock interacting clouds using the CO/H I radio line
emission. Most recently, Sano et al. (2021) presented the first reliable
results of an SNR age–$W_{\mathrm{p}}$ relation for eleven gamma-ray SNRs. The
authors discovered a positive correlation between the SNR age and
$W_{\mathrm{p}}$ with ages below $\sim 6$ kyr, suggesting that in-situ values
of $W_{\mathrm{p}}$ for young SNRs are limited by the short duration time of
cosmic-ray acceleration, also known as the age-limited acceleration (cf. Ohira
et al., 2010). The older SNRs with ages more than $\sim 8$ kyr, on the other
hand, show a steady decrease of $W_{\mathrm{p}}$. This can be explained as an
effect of the energy-dependent diffusion of cosmic rays (e.g., Aharonian &
Atoyan, 1996; Gabici et al., 2007). In this section, we argue that
G346.6$-$0.2 can be also naturally understood by the paradigm of the SNR
age–$W_{\mathrm{p}}$ relation.

To derive the $W_{\mathrm{p}}$ value of G346.6$-$0.2, we first estimate the
masses and number densities for the shock-interacting molecular and atomic
clouds. The mass of molecular clouds $M_{\mathrm{CO}}$ can be estimated using
the following equations:

$\displaystyle M_{\mathrm{CO}}=m_{\mathrm{p}}\mu\Omega
D^{2}\sum_{i}N_{i}(\mathrm{H}_{2}),$ (2) $\displaystyle
N(\mathrm{H}_{2})=X\cdot W(\mathrm{CO}),$ (3)

where $m_{\mathrm{p}}$ is the mass of atomic hydrogen, $\mu=2.8$ is the mean
molecular weight, $\Omega$ is the solid angle for each pixel, $D$ is the
distance to the SNR, $N(\mathrm{H}_{2})$ is the molecular hydrogen column
density in units of cm-2, $X$ is the CO-to-H2 conversion factor of $2\times
10^{20}$ cm-2 (K km s${}^{-1})^{-1}$ (Bertsch et al., 1993), and $W$(CO) is
the velocity integrated intensity of 12CO($J$ = 1–0) emission line in units of
K km s-1. We estimated the mass of molecular clouds within the radio shell
extent ($=$ shell radius $+$ $1/2$ shell thickness) to be $\sim 4.1\times
10^{4}$ $M_{\sun}$. The mass of atomic hydrogen clouds $M_{\mathrm{HI}}$ can
be also given by the equations (4) and (5) under the optically thin assumption
(Dickey & Lockman, 1990):

$\displaystyle M_{\mathrm{HI}}=m_{\mathrm{p}}\Omega
D^{2}\sum_{i}N_{i}(\mathrm{H{\textsc{i}}}),$ (4) $\displaystyle
N(\mathrm{H{\textsc{i}}})=1.823\times W(\mathrm{H{\textsc{i}}}),$ (5)

where $N$(H I) is the atomic hydrogen column density in units of cm-2 and
$W$(H I) is the velocity integrated intensity of H I line emission in units of
K km s-1. We then estimated the mass of atomic clouds to be $\sim 0.5\times
10^{4}$ $M_{\sun}$, which is eight times less than the mass of molecular
clouds. This trend is roughly consistent with other middle-aged SNRs (e.g.,
Yoshiike et al., 2013, 2021; Kuriki et al., 2018). The averaged number
densities within the wind-shell could be estimated to $\sim 120$ cm-3 for the
molecular hydrogen $n$(H2), and $\sim 40$ cm-3 for the atomic hydrogen $n$(H
I) by adopting the shell radius of 3$\farcm$29 ($\sim 10.6$ pc) and the shell
thickness of 1$\farcm$41 ($\sim 4.6$ pc). Then we derived the number density
of the total interstellar protons $n$ to be 2 $\times$ $n$(H2) $+$ $n$(H I) =
280 cm-3.

The total energy of accelerated cosmic-rays $W_{\mathrm{p}}$ can be derived by
using the following equation if the hadronic process is dominantly working
(e.g., Aharonian et al., 2006):

$\displaystyle W_{\mathrm{p}}\sim t_{\mathrm{pp\rightarrow\pi^{0}}}\times
L_{\gamma}$ (6)

where $t_{\mathrm{pp\rightarrow\pi^{0}}}\sim 4.5\times 10^{15}$ ($n$/1 cm-3)-1
s is the characteristic cooling time of cosmic-ray proton and $L_{\gamma}$ is
the gamma-ray luminosity in units of erg s-1. By adopting $L_{\gamma}$(1–100
GeV$)<5.8\times 10^{34}$ erg s-1 and $n=280$ cm-3 at the SNR distance of 11.1
kpc, we finally derived $W_{\mathrm{p}}$(10–1000 GeV$)<9.3\times 10^{47}$ erg,
corresponding to less than $\sim 0.1$% of the typical released kinetic energy
of $\sim 10^{51}$ erg of a supernova explosion.

Figure 7 shows the SNR age–$W_{\mathrm{p}}$ diagram of twelve gamma-ray SNRs
including G346.6$-$0.2. We can clearly see that G346.6$-$0.2 shows roughly
consistent values with the other SNRs such as IC 443 (Yoshiike et al., 2021).
Such a low energy amount suggests that most of the cosmic-rays in G346.6$-$0.2
have been escaped from the SNR, we can find the small in-situ value of
$W_{\mathrm{p}}$. Further gamma-ray observations will possibly detect the
hadronic gamma-ray emission due to the escaped cosmic-rays in the vicinity of
the SNR.

##  5 Conclusions

We summarize the primary conclusions as follows:

  1. 1.

We revealed CO and H I clouds associated with the mixed-morphology SNR
G346.6$-$0.2 using the NANTEN2 12CO($J$ = 2–1) and ATCA & Parkes H I datasets.
The H I clouds are distributed surrounding the radio continuum shell except
for the southeastern direction, while the CO clouds are located only the south
and northeast shell of the SNR. The south CO clouds show a good spatial
correspondence with the positions of the shock-excited 1720 MHz OH masers at
$V_{\mathrm{LSR}}\sim-70$ km s-1. The northeast CO cloud shows a high-kinetic
temperature of $\sim 60$ K, suggesting that the shock heating occurred.

  2. 2.

The cavity-like structure in the position-velocity diagram of H I (possibly CO
as well) indicates a wind-blown bubble which was likely formed by strong
stellar winds from the high-mass progenitor of the SNR. The systemic velocity
and expansion velocity of the wind bubble were derived to $\sim-72$ km s-1 and
$\sim 10$ km s-1, respectively. The H I absorption feature due to the SNR is
seen in the velocity range from $-95$ to $-2$ km s-1, suggesting that
G346.6$-$0.2 is located on the far side of the Galactic center from us and the
appropriate distance is to be $11.1^{+0.5}_{-0.3}$ kpc. The plasma age of the
SNR was also revised to be $14\pm 2$ kyr.

  3. 3.

We found that the X-ray recombining plasma shows a clear spatial anti-
correlation with the shocked CO clouds. In addition, there is no specific
correlation between the CO intensities and the electron temperature of the
recombining plasma. These results favor the adiabatic cooling scenario as the
formation mechanism of recombining plasma in G346.6$-$0.2, rather than the
thermal conduction scenario. This is also consistent with the presence of the
stellar wind bubble because the dense CSM will be expected in the vicinity of
the high-mass progenitor.

  4. 4.

Using the latest Fermi-LAT datasets, we placed a conservative upper limit of
gamma-ray luminosity from the SNR G346.6$-$0.2 to $\sim 5.8\times 10^{34}$ erg
s-1 at the distance of 11.1 kpc. This corresponds to the total energy of
cosmic-rays $W_{\mathrm{p}}<9.3\times 10^{47}$ erg adopting the target
interstellar gas density of 280 cm-3 if the hadronic process is dominantly
working. The SNR age–$W_{\mathrm{p}}$ relation indicates that most of the
accelerated cosmic-rays have been escaped from the SNR shell.

##

The NANTEN project is based on a mutual agreement between Nagoya University
and the Carnegie Institution of Washington (CIW). We greatly appreciate the
hospitality of all the staff members of the Las Campanas Observatory of CIW.
We are thankful to many Japanese public donors and companies who contributed
to the realization of the project. NANTEN2 is an international collaboration
of 10 universities, Nagoya University, Osaka Prefecture University, University
of Cologne, University of Bonn, Seoul National University, University of
Chile, University of New South Wales, Macquarie University, University of
Sydney, and ETH Zurich. The Mopra telescope, Australia Telescope Compact Array
(ATCA), and the Parkes radio telescope are parts of the Australia Telescope
National Facility which is funded by the Australian Government for operation
as a National Facility managed by CSIRO. We acknowledge the Gamilaroi,
Gomeroi, and Wiradjuri people as the traditional owners of the Observatory
sites. The University of New South Wales Digital Filter Bank used for the
observations with the Mopra Telescope was provided with support from the
Australian Research Council. The Molonglo Observatory Synthesis Telescope
(MOST) is operated by the University of Sydney with support from the
Australian Research Council and the Science Foundation for Physics within the
University of Sydney. This paper made use of information from the SEDIGISM
survey database located at https://sedigism.mpifr-bonn.mpg.de/index.html,
which was constructed by James Urquhart and hosted by the Max Planck Institute
for Radio Astronomy. Based on observations obtained with XMM-Newton, an ESA
science mission with instruments and contributions directly funded by ESA
Member States and NASA. This work was supported by JSPS KAKENHI Grant Numbers
JP19H05075 (H. Sano), JP20K14491 (K. Nobukawa), JP20KK0071 (K. Nobukawa),
JP21H01136 (H. Sano), JP21K03615 (K. Nobukawa), and JP21J00031 (H. Suzuki). K.
Nobukawa was also supported by Yamada Science Foundation.

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