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arXiv:1102.5333v1 [astro-ph.HE] 25 Feb 2011
ACCEPTED TO APJ: FEBRUARY 14, 2011
Preprint typeset using L
A
TEX style emulateapj v. 11/10/09
A CHANDRA OBSERVATION OF SUPERNOVA REMNANT G350.1−0.3
AND ITS CENTRAL COMPACT OBJECT
I. LOVC HINSKY1, P. SL ANE1, B. M. GAENSLER2, J. P. HUGHES3, C.-Y. NG4, J. S. LAZE NDIC5, J. D. GEL FAND6, & C. L. BROGAN7
ACCEPTED TO APJ: February 14, 2011
ABSTRACT
We present a new Chandra observation of supernova remnant (SNR) G350.1−0.3. The high resolution X-
ray data reveal previously unresolved filamentary structures and allow us to perform detailed spectroscopy
in the diffuse regions of this SNR. Spectral analysis demonstrates that the region of brightest emission is
dominated by hot, metal-rich ejecta while the ambient material along the perimeter of the ejecta region and
throughout the remnant’s western half is mostly low-temperature, shocked interstellar/circumstellar medium
(ISM/CSM) with solar-type composition. The data reveal that the emission extends far to the west of the
ejecta region and imply a lower limit of 6.6 pc on the diameter of the source (at a distance of 4.5 kpc). We
show that G350.1−0.3 is likely in the free expansion (ejecta-dominated) stage and calculate an age of 600 −
1200 years. The derived relationship between the shock velocity and the electron/proton temperature ratio is
found to be entirely consistent with that of other SNRs. We perform spectral fits on the X-ray source XMMU
J172054.5−372652, a candidate central compact object (CCO), and find that its spectral properties fall within
the typical range of other CCOs. We also present archival 24 µm data of G350.1−0.3 taken with the Spitzer
Space Telescope during the MIPSGAL galactic survey and find that the infrared and X-ray morphologies are
well-correlated. These results help to explain this remnant’s peculiar asymmetries and shed new light on its
dynamics and evolution.
Subject headings: ISM: individual (G350.1−0.3) — stars: individual (XMMU J172054.5−372652) — stars:
neutron — supernova remnants
1. INTRODUCTION
Observations of young supernova remnants (SNRs) provide
our best means of relating the thermodynamic properties of
shock waves and ejecta distributions to the mechanics of su-
pernovae and their associated progenitors. Chandra studies
of young remnants have yielded some of the most spectacular
advances in SNR research and revealed intriguing and unex-
pected features of these intricate systems.
G350.1−0.3 is a young, luminous SNR in the inner Galaxy.
Its nonthermal and linearly polarized radio emission led to its
identification as a SNR by Clark et al. (1973,1975), but a
4.8 GHz Very Large Array radio image (Salter et al. 1986)
revealed a bizarre, asymmetric morphology, unlike that of
the typical shell-like structures. In subsequent SNR catalogs,
the source was either removed or listed as a SNR candidate
(Green 1991; Whiteoak & Green 1996). G350.1−0.3 has been
detected with ROSAT and ASCA, where it was labeled 1RXS
J172106.9−372639 and AX J1721.0−3726, respectively (Vo-
ges et al. 1999, Sugizaki et al. 2001). An observation with
XMM-Newton (Gaensler et al. 2008, hereafter G08) revealed
1Harvard-Smithsonian Center for Astrophysics, 60 Garden
Street, Cambridge, MA 02138, USA; ilovchin@fas.harvard.edu;
slane@cfa.harvard.edu
2Sydney Institute for Astronomy, School of Physics, The University of
Sydney, NSW 2006, Australia; bryan.gaensler@sydney.edu.au
3Rutgers University, The State University of New Jersey, Piscataway,
NJ, USA; jph@physics.rutgers.edu
4Department of Physics, McGill University, Montreal, QC H3A 2T8,
Canada; ncy@hep.physics.mcgill.ca
5School of Physics, Monash University Clayton, VIC 3800, Australia;
Jasmina.Lazendic-Galloway@monash.edu
6New York University Abu Dhabi, PO Box 129188, Abu Dhabi, United
Arab Emirates; jg168@astro.physics.nyu.edu
7National Radio Astronomy Observatory, 520 Edgemont Rd, Char-
lottesville VA, 22903, USA; cbrogan@cv.nrao.edu
the source to be concentrated in a bright clump, coincident
with the region of strongest radio emission. The X-ray spec-
trum was found to be well-fit by a two-component model
(XSPEC model VPSHOCK + RAYMOND), consisting of a
shocked plasma with electron temperature kTe≈1.5 keV, ion-
ization timescale τ≈3×1011 s cm−3and large overabun-
dances of all metals, as well as a collisionally equilibrated,
low-temperature (kTe≈0.4 keV) component with solar abun-
dances. The best-fit value for interstellar absorption was
NH≈3.8×1022 cm−2. G08 proposed that G350.1−0.3 is in-
teracting with a molecular cloud seen along its eastern edge
in the 12CO survey of Bitran et al. (1997). These considera-
tions as well as data from the Southern Galactic Plane Survey
(McClure-Griffiths et al. 2005) were used to derive an ap-
proximate distance of 4.5 kpc. Henceforth we scale quantities
with d4.5, the distance in units of 4.5 kpc.
G08 detected an unresolved X-ray source, XMMU
J172054.5−372652, to the west of the bright emission and
proposed it to be a neutron star, likely a central compact object
(CCO), associated with G350.1−0.3. Its spectrum was found
to be a good fit to an absorbed blackbody although a power-
law model with an unphysically high photon index (Γ= 5.4)
provided an acceptable fit as well. G08 carried out a search
for possible pulsations but found no periodic signals in the
range of 146 ms to 1 hr.
In this paper, we present a detailed study of G350.1−0.3
and its candidate CCO. In § 2 we summarize the Chandra and
Spitzer observations and data reduction. § 3.1 features a dis-
cussion of the X-ray and infrared imaging and a comparison
of various spectral fits. In § 3.2 we use the parameters de-
rived from the models to analyze the dynamics and evolution
of G350.1−0.3. A discussion of the compact source XMMU
J172054.5−372652 follows in § 3.3. In § 4 we interpret the
results of these investigations and attempt to present a coher-
2 Lovchinsky et al.
15.0 10.0 05.0 17:21:00.0 55.0 20:50.0
24:00.0
25:00.0
26:00.0
-37:27:00.0
28:00.0
29:00.0
0.0
0.2
0.5
1.1
2.3
4.7
9.5
19.1
38.5
76.9
153.2
11.0 17:21:10.0 09.0 08.0 07.0 06.0 05.0 04.0
26:20.0
30.0
40.0
50.0
-37:27:00.0
10.0
20.0
30.0
40.0
50.0
Right ascension (J2000)
Declination (J2000)
0
0.042
0.12
0.29
0.62
1.3
2.6
5.3
11
21
42
2
3
4
1a
1b
1c
(a)
(b)
Figure 1. Panel ashows a Chandra image of G350.1−0.3 in the 0.5−8.0 keV
band. The image has been smoothed by a Gaussian kernel with a σof 3 pixels
(1.48 arcsec). The detailed structure of the bright emission region in the east
is not visible due to the high contrast. The dashed contour surrounding the
SNR is at a level of 3σabove the background and shows the apparent extent
of the emission. The white cross indicates the centroid of the contour and is
regarded as the putative center of the SNR throughout this analysis. Regions
2−4 were used for spectral extraction (as shown in Figure 3) with parameters
listed in Table 1. The dashed circle, partly shown in the southwest corner,
was used to subtract the local background. The arrow indicates the position
of the candidate CCO XMMU J172054.5−372652. Panel bshows an enlarged
image of the bright eastern emission indicated by the square region in panel a.
Regions 1a, 1b and 1c were used for spectral extraction (as shown in Figure
2), with parameters for regions 1a and 1b listed in Table 1. The colorbars (in
units of counts/pixel) indicate brightness levels.
ent and self-consistent picture of G350.1−0.3. We conclude
with § 5, a summary of our findings.
2. OBSERVATIONS
In order to obtain a high-resolution map of the SNR emis-
sion structure, study its spectral properties in unprecedented
detail and obtain tighter constraints on the spectral fit param-
eters, we observed G350.1−0.3 on 2009 May 21 and 22 for
84 ks with the Advanced CCD Imaging Spectrometer (ACIS-
S) on board the Chandra X-ray Observatory (observation ID
10102). Standard data reduction was performed to remove
hot pixels and flares, resulting in an effective exposure time
of 83 ks. The 24 µm and 70 µm images of the spatial
region corresponding to G350.1−0.3 were obtained during
MIPSGAL, an infrared survey of the Galactic plane (Carey
et al. 2009). The target was observed (in both bands) on
2006 October 06 with the Multiband Infrared Photometer for
Spitzer (MIPS) instrument on board the Spitzer Space Tele-
scope. The raw data were calibrated using the standard BCD
(basic calibrated data) pipeline to correct for bright pixels
and flares resulting in exposure times of 2.6 s (24 µm) and
3.15 s (70 µm). The 70 µm data were additionally filtered
to mitigate instrumental effects (stimflash latents and resid-
ual background drifts). An 8 µm image was obtained during
GLIMPSE, an infrared survey of the inner Milky Way Galaxy
(Churchwell et al. 2009). The target was observed on 2004
September 07 with the InfraRed Array Camera (IRAC) on
board the Spitzer Space Telescope. The raw data were cali-
brated using the standard BCD pipeline to correct for hot pix-
els and flares resulting in an exposure time of 1.2 s.
3. ANALYSIS
3.1. Imaging and Spectroscopy
AChandra image of G350.1−0.3 is shown in Figure 1a.
We filtered the data on the 0.5 - 8.0 keV energy band and
smoothed the image by applying a Gaussian kernel with a σ
of 3 pixels (1.48 arcsec). The source is dominated by an ex-
tended bright clump of material in the east with fainter emis-
sion extending far to the west. The dashed contour, shown
in the figure, surrounds the SNR at a level of 3σabove the
background and indicates the apparent extent of the emis-
sion. Throughout our analysis, we adopt a source radius of
R= 3.3d4.5pc (2.5 arcmin) and assume, following the argu-
ment of G08, that d4.5= 1. Although the data reveal an el-
liptical shell-like structure, a large region in the southwest is
markedly fainter than the adjacent emission with a flux only
a factor of ∼1.5 above the background. The unresolved X-
ray source XMMU J172054.5−372652 (indicated by the ar-
row in Figure 1a) is seen far to the west of the bright emis-
sion, significantly displaced from the apparent center of the
SNR (marked by the white cross in Figure 1a). Figure 1b
shows the detailed structure of the ejecta region and reveals
previously unresolved filamentary structures running across
its center. We find that the six regions enumerated in Fig-
ures 1a and 1b represent the range of spectral characteris-
tics found in this remnant. The spectra and the correspond-
ing response files were produced using the specextract script
in Ciao 4.2. We subtracted the local background spectrum
from a source-free region of the detector, partly shown as
the dashed region in the southwest corner of Figure 1a, and
grouped the data to a minimum of 15 counts per bin. We
fit each spectrum by an absorbed nonequilibrium ionization
(NEI) plane-parallel shock model with variable abundances
(XSPEC model "VNEI"; Borkowski et al. 2001). Elements
below Mg and above Fe were fixed at solar abundances since
their contributions in the fitted bandpass are negligible. The
uncertainties were calculated at 90% confidence (1.6σ). Al-
though fixing the absorption at an average global value of
∼4.0×1022 cm−2produced acceptable fits for regions 1a, 1b,
2, 3 and 4 (the exception is region 1c and is discussed be-
low), we allow NHto vary in order to account for possible
small-scale variations in the absorbing column. The spectra
and fitted models for regions 1a, 1b and 1c are shown in Fig-
ure 2 with parameters (for regions 1a and 1b) listed in Table
1. The fits for regions 2, 3 and 4 are shown in Figure 3, with
parameters listed in Table 1.
Chandra Observation of G350.1−0.3 3
0
10
10
-0
10
-1
10
-2
10
-
3
-5
-5
0
5
10
-1
10
-2
10
-3
10
-
4
counts s -1 keV -1
counts s -1 keV -1
counts s -1 keV -1
0
5
10
-1
10
-2
10
-3
10
-
4
-10
1 2 3 4 5 6 7
energy (keV)
(1a)
(1b)
(1c)
Mg
Si
S
Ar
Ca
Fe
Mg
Si
S
Ar
Ca
Fe
Mg Si S
Ar
Ca
Fe
Figure 2. Plots of spectra extracted from regions 1a, 1b and 1c overlaid
with the fitted VNEI models and residuals. The NEI model in XSPEC does
not include atomic data for the argon emission; a Gaussian component was
added to the model to account for this feature in the spectral fitting. The fit
parameters for regions 1a and 1b are listed in Table 1.
The spectrum of region 1a (extracted from one of the
bright filaments in Figure 1b) is roughly representative of
that found throughout the bright eastern emission. It can
be fit by an absorbed VNEI model with column density
NH≈3.8×1022 cm−2, electron temperature kTe≈1.4 keV,
ionization timescale τ≈2.2×1011 s cm−3and large over-
abundances of Mg, Si, S, Ca and Fe. These characteristics
unambiguously demonstrate the presence of hot, metal-rich
ejecta. The spectrum of region 1b (extracted from the other
bright filament) is similar to that of region 1a although the
Mg and Si lines clearly indicate a somewhat lower ionization
timescale. Although morphologically well-defined, we find
that the spectral characteristics of the filaments are basically
consistent with the surrounding regions.
Although G08 found evidence for variations in elemental
abundances between nearby regions within the bright east-
ern emission, our spectral fits of small regions in this vicinity
show no significant variations within uncertainties. In order to
compare our results with those of the XMM-Newton observa-
tion, we attempted to fit the entire ejecta-dominated emission
(region 1c: the same area as fitted by G08) using a VNEI
model. As can be seen from the high residuals, this model
does not provide a good fit, and hence, we do not display the
1 2 3 4 5 6 7
energy (keV)
-5
0
5
10
-1
10
-2
10
-3
10
-
4
(3)
(2)
(4)
Mg
Si
S
Ar
Ca
Mg
Si
S
Ar Ca
Mg
Si
S
Ar
Ca
-5
0
5
10
-1
10
-2
10
-3
10
-
4
-5
0
5
10
-1
10
-2
10
-3
10
-
4
counts s -1 keV -1
counts s -1 keV -1
counts s -1 keV -1
Figure 3. Plots of spectra extracted from regions 2, 3 and 4 overlaid with
the fitted VNEI models and residuals. The fit parameters are listed in Table
1.
fit parameters. However, the best-fit parameters are roughly
consistent with those of regions 1a and 1b, indicating that the
extremely large number of X-ray photon counts (∼100,000)
and small error bars likely magnify small deficiencies in the
model and result in the poor χ2statistics. In addition, it is
possible that small variations in temperature and elemental
abundances exist, but are not discernible with the statistics in
our spectra from small-scale, spatially-resolved regions. Us-
ing the two-component model VPSHOCK + RAYMOND (the
same as that used by G08 - their fit results are summarized in
§ 1) for region 1c also fails to provide a good fit. The incon-
sistencies between our results and those of G08 are likely ex-
plained simply by the fact that the long exposure in the Chan-
dra observation exposes weaknesses in the models that were
previously hidden due to the fact that the XMM-Newton spec-
trum has fewer photon counts.
Region 3 is largely representative of the spectral features
found along the perimeter of the ejecta region and throughout
the remnant’s western half. The spectrum can be described by
a VNEI model with absorbing column NH≈4.4×1022 cm−2,
electron temperature kTe≈0.85 keV, ionization timescale
τ≈2.7×1011 s cm−3and near-solar abundances of all met-
als. These results suggest an interstellar/circumstellar origin.
Although otherwise similar, the fit for region 4 indicates a
somewhat lower temperature and an ionization timescale of
4 Lovchinsky et al.
Table 1
VNEI parameters for spectral fits of regions 1a, 1b, 2, 3, 4
Parameter Region 1a Region 1b Region 2 Region 3 Region 4
NHa(1022 cm−2) 3.8+0.4
−0.24.0+1.3
−0.13.9+0.3
−0.24.4+0.3
−0.24.3+0.2
−0.2
kTe(keV) 1.4+0.2
−0.21.2+0.2
−0.10.8+0.1
−0.10.9+0.1
−0.10.54+0.04
−0.04
Mg 5.8+8.2
−1.88.6+1.9
−1.82.6+1.1
−0.72.2+0.9
−0.61.1+0.2
−0.2
Si 8.6+9.3
−1.87.2+1.4
−1.42.8+0.5
−0.61.6+0.3
−0.21.7+0.3
−0.1
S 4.4+4.4
−0.93.0+0.6
−0.62.3+0.4
−0.41.2+0.2
−0.20.5+0.2
−0.1
Ca 6.1+2.0
−1.84.9+2.3
−2.36.3+3.1
−2.41.1+1.0
−0.90+4.9
−−
Feb4.3+8.1
−1.64.3+1.4
−1.4(1) (1) (1)
τ(1011 s cm−3) 2.2+0.6
−0.30.92+0.10
−0.01 4.5+2.4
−1.22.7+0.5
−0.50.13+0.04
−0.02
Fluxc(10−13 ergs cm−2s−1) 8.9 8.2 5.4 11 7.2
χ2
ν/dof 1.94/157 2.16/144 1.89/115 1.55/178 2.27/154
Note. — All uncertainties are statistical errors at 90% confidence, 1.6σ. Elements below Mg
and above Fe were fixed at solar because their contributions in the fitted bandpass are small
aThe absorption was calculated using the model of Wilms et al. (2000)
b(1) indicates that the elemental abundance was fixed at solar
cThe values listed are the unabsorbed fluxes over the energy range 0.5−8.0 keV
τ≈1.3×1010 s cm−3, lower than that of region 3 by a factor
of 20. The centroid of the Si line in the region 4 spectrum
is at a slightly lower energy than that of region 3 and since
the goodness of fit statistics are dominated by the Si line (the
error bars are smallest in this region), it is this feature that
causes the drastic difference in the ionization timescale. Re-
gion 2, adjacent to the bright eastern emission, has somewhat
enhanced abundances of Mg, Si and S and a highly enhanced
abundance of Ca. Although similar in temperature to region
3, the overabundance of metals suggests the possible presence
of ejecta.
The model histogram for the spectrum of region 4 slightly
misses the centroid of the S line complex, suggesting that
the several bright patches within this region likely vary in
their ionization states. Fitting the individual clumps, how-
ever,leads to a similar result with the added handicap of poor
statistics in the spectra. We thus present the fit to the com-
posite region with the understanding that a deeper exposure
is necessary to characterize the spectrum more precisely. The
spectral fits to regions 1b, 2 and 4 do not formally correspond
to good fits based on the χ2statistics. In the case of region
1b, this deficiency is due to the fact that the model histogram
misses the minimum of the Mg line and is too low on the low-
energy shoulder of the Si line. In the region 2 spectrum, the
fit is too low on the maximum of the Si line, as can be seen
from the high residuals in this area. We acknowledge these
weaknesses in the models and use them to simply make qual-
itative statements about the approximate distribution of ele-
mental abundances, temperatures and ionization states, rather
than using them to derive precise numerical results.
A 24 µm infrared image of G350.1−0.3, overlaid with the
contours from the X-ray data, is shown in Figure 4a. The
two bright clumps in the east and the south correspond to re-
gions 1c and 4 respectively. Most notably, the upper filament
from the X-ray image (region 1b) appears to outline precisely
the eastern edge of the infrared emission while the lower fil-
ament (region 1a) is notably absent at 24 µm. An enlarged
24 um image of the ejecta- dominated region, overlaid with
the X-ray contours, is shown in Figure 4b. While the angu-
lar resolution at 24 um is much coarser than that provided
by Chandra, these images are sufficiently sensitive to iden-
tify two distinct regions of emission with clear morphological
similarities to the X-ray emission from the SNR (regions 1c
and 4 in Figure 1). Figures 4c and 4d show IRAC 8 µm and
MIPS 70 µm images of the spatial region corresponding to
G350.1−0.3. We do not detect any significant emission from
G350.1−0.3 at 8 µm and although we see 70 µm emission
in regions corresponding to the SNR, the emission does not
differ significantly from regions in the surrounding medium.
With the current resolution we cannot determine whether or
not some of the 70 µm emission is associated with the SNR.
We searched the 3.6 µm, 4.5 µm and 5.8 µm IRAC bands
but found no apparent emission from G350.1−0.3 or XMMU
J172054.5−372652.
3.2. Dynamics and Evolution
We analyze the dynamics and evolution of G350.1−0.3 us-
ing the numerical study of nonradiative SNRs by Truelove
and McKee (1999), hereafter TM99. We fit the X-ray spec-
trum of the SNR’s entire western half by an absorbed VNEI
model in order to derive the average post-shock electron den-
sity and the mass of the swept-up material. The parameters
of the model are roughly intermediate to those of regions 3
and 4. We assume a spherical half-shell with thickness R/12
for the geometry and a value of 4 for the compression ratio
of post-shock and pre-shock densities. We also take the ra-
tio between the electron and atomic hydrogen densities to be
ne/nH= 1.2, valid for cosmic abundances. With the normal-
ization parameter from the fit (∼0.056), we use the relation
nenHV
4πd2= 0.056×1014 cm−5,(1)
where Vis the estimated volume and dis the distance,
to derive an approximate post-shock electron density ne≈
5.6d−1/2
4.5cm−3and a corresponding swept-up mass of ∼
2d−3/2
4.5M⊙, with the obvious caveat that the uncertainties are
unquantifiable since the visible emission — and hence the
observed radius — is only a lower constraint on the true
size of the SNR. G08 derived a substantially higher density
through two independent methods. However, the first calcu-
lation assumes ionization equilibrium, inferred from using a
Chandra Observation of G350.1−0.3 5
10.0 05.0 17:21:00.0 55.0 20:50.0
25:00.0
30.0
26:00.0
30.0
-37:27:00.0
30.0
28:00.0
30.0
Declination (J2000)
74.8
75.4
76.6
79.0
83.8
93.4
112.4
150.3
226.7
377.8
678.7
(a)
1c
4
10.0 05.0 17:21:00.0 55.0 20:50.0
25:00.0
30.0
26:00.0
30.0
-37:27:00.0
30.0
28:00.0
30.0
13
19
32
57
106
206
404
797
1592
3163
6291
(c)
Declination (J2000)
1c
4
10.0 05.0 17:21:00.0 55.0 20:50.0
25:00.0
30.0
26:00.0
30.0
-37:27:00.0
30.0
28:00.0
30.0
Right ascension (J2000)
-80
-66
-38
19
131
358
807
1699
3501
7066
14163
(d)
Declination (J2000)
1c
4
Right ascension (J2000)
11.0 17:21:10.0 09.0 08.0 07.0 06.0 05.0 04.0
26:20.0
30.0
40.0
50.0
-37:27:00.0
10.0
20.0
30.0
40.0
50.0
03.0 02.0
74.8
75.4
76.6
79.0
83.8
93.4
112.4
150.3
226.7
377.8
678.7
Declination (J2000)
(b)
Figure 4. Panel ashows a 24 µm image of G350.1−0.3 taken during the MIPSGAL Galactic survey. The integrated flux of regions 1c and 4 is ∼2.5 Jy and
∼0.5 Jy respectively. Panel bshows an enlarged image ofthe region corresponding to the ejecta-dominated emission with corresponding X-ray contours at levels
of 1.1, 2.5, 7.0 and 13 counts/pixel. Panels cand dshow 8 µm and 70 µm images of the same spatial region as panel a. The contours in panels a,cand dare
at levels of 0.55, 1.1, 5.0 and 13 counts/pixels and correspond to the X-ray data from Figure 1a. Regions 1c and 4 correspond to the same regions as in Figure 1.
The colorbars indicate brightness levels in units of MJy/sr.
two-component model (which we were unable to fit, as dis-
cussed in § 3.1). The second calculation assumes a Sedov
solution — an assumption that is likely not valid for reasons
discussed below. If the molecular cloud interaction scenario
(as proposed by G08) is correct, the density throughout the
ejecta-dominated half of the SNR is likely to be higher than
that in the western half. Hence, extrapolating our estimate to
the eastern half of the SNR (for a total of ∼4d−3/2
4.5M⊙) is
likely an underestimate of the true swept-up mass, although
this figure can be used as a rough lower bound.
For a progenitor with an ejecta density profile ∝r−7(typical
of core-collapse supernovae), TM99 demonstrate that the ra-
dius and age at which a free-expanding SNR enters the Sedov-
Taylor (ST) phase are given by
RST = 0.881M1/3
e j ρ−1/3
0(2)
and tST = 0.732E−1/2
SN M5/6
e j ρ−1/3
0,(3)
where Me j is the ejecta mass and ESN is the explosion energy.
Here ρ0is the pre-shock mass density and is given by
ρ0=1.4
(1.2)(4)mHne= 0.30mHne,(4)
where mHis the mass of the hydrogen atom, 1.4 is the equiv-
alent molecular weight of the hydrogen and helium mix-
ture, (assuming cosmic abundances) and the factors 1.2 and
4 are the electron/hydrogen density ratio and post-shock/pre-
shock compression ratio, respectively. Using typical values
of Me j = 4 −14 M⊙for the ejecta mass consistent with the
formation of a neutron star (Woosley et al. 2002), we calcu-
late RST to be 4.1−6.2 pc, somewhat larger than the observed
radius for G350.1−0.3 (R= 3.3d4.5pc). Assuming that the ob-
served radius of the SNR is approximately equal to its true
radius (in § 4 we discuss the scenario in which this assump-
tion is not valid), this result suggests that G350.1−0.3 is likely
in the free expansion (ejecta-dominated) stage and has not yet
entered the ST phase.
If we make an assumption about the explosion energy, we
can determine the age using the relation
t= 0.988tST R
RST 7/4
,(5)
as derived from TM99. We explore the parameter space for
Me j = 4 −14 M⊙and ESN = 0.5−1.0×1051 ergs and derive an
age of t= 600 −1200 years. The shock velocity and proton
6 Lovchinsky et al.
Figure 5. Azimuthally-averaged radial profile of CCO XMMU
J172054.5−372652 in the 0.5−8.0 keV energy range. The histogram with
the vertical error bars represents the data while the curve is the model point
spread function (PSF).
temperature follow from8
vb=dR
dt = 0.576RST
tST t
tST −3/7(6)
and kTp= 0.11mHv2
b.(7)
The results of these calculations are summarized in Table
2. Using kTe≈0.7 keV for the electron temperature of the
ISM component, the average value from our fits, the inferred
electron-to-proton temperature ratio is 0.07 −0.27. This is
in excellent agreement with measured values of this ratio for
SNR shocks in the inferred velocity range (Ghavamian et al.
2007).
3.3. The Compact Source
Using the derived age for the remnant, we calculate the pro-
jected velocity of XMMU J172054.5−372652 to be 1400 −
2600 km s−1, assuming a displacement of 1.7d4.5pc (1.3 ar-
cmin) from the adopted center of the SNR. We generated a
lightcurve of XMMU J172054.5−372652 but found no short-
term flux variation within the Chandra exposure. We ap-
plied barycenter corrections to the photon arrival times and
searched for periodicity using the Z2-test (Buccheri et al.
1983). No pulsations with period longer than 6.4 s and pulsed
fraction larger than 16% were detected at 99% confidence.
The Chandra image in Figure 1a shows no extended emission
surrounding the CCO and a generated radial profile (Figure 5)
of the source confirms that it is fully consistent with a model
PSF. As a note, photon pile-up is negligible with a level of
less than 5%.
8This expression for the shock velocity comes from TM99 Table 7. It
should be noted, however, that the authors overlooked the minus sign in the
exponent.
Table 2
Derived Parameters for the Evolution of G350.1−0.3
Parameter Value
Age, t, (years) 600−1200
Shock Velocity, vb, (kms−1) 1500−2900
Proton Temperature, kTp, (keV) 3−10
Speed of CCO, vcco, (km s−1) 1400−2600
We extracted the ∼4000 CCO counts from a 2 arcsec-
radius aperture and grouped the source spectrum to a mini-
mum of 25 counts per bin. Using the SHERPA environment,
we fit the spectrum to simple absorbed blackbody (BB) and
power-law (PL) models. Although the PL fit is slightly bet-
ter than the BB, the large photon index (Γ= 5.5) suggests a
thermal origin for the emission. Both fits are consistent with
those discussed in G08. Since CCOs are often characterized
by two-component models (BB+PL and BB+BB), we inves-
tigated those as well. The results are summarized in Table
3. The best-fit absorption column density is ∼4×1022 cm2
and agrees with that of the SNR. While the CCO spectrum
is dominated by a BB with kT ≈0.4 keV and a small emis-
sion radius of ∼3 km, the high energy component is not very
well constrained. Our results suggest that it can be modeled
by either an additional BB component with a very high tem-
perature (0.9 keV), or a PL component. Adding a second
component improves the fit although not at a statistically sig-
nificant level. Unfortunately, the current data do not allow
us to distinguish between these two scenarios and a deeper
exposure is required. We present the spectrum of XMMU
J172054.5−372652 overlaid with the BB+PL model in Figure
6.
4. DISCUSSION
The high resolution Chandra data reveal that the emission
from G350.1−0.3 extends far to the west of the bright ejecta
region and provide a rough constraint on the size of the SNR,
with the caveat that the asymmetric morphology makes it dif-
ficult to extract a reliable diameter. G08 estimated the age of
G350.1−0.3 to be ∼900 years based on the diameter of the
ejecta region alone (D≈2.6d4.5pc) and the shock velocity
derived from the electron temperature. However, we see that
the emission from G350.1−0.3 extends far beyond the bright
ejecta-dominated region and our analysis (using the model
treatment of TM99) shows that the electron temperature is
almost certainly an underestimate of the proton temperature.
With the diameter, a post-shock density estimate and an as-
sumed range of values for Me j and ESN as input parameters,
we calculate a shock velocity of 1500 −2900 km s−1, a proton
temperature of 3 −10 keV and an age of 600−1200 years. We
see that although the size of the SNR is far larger, the age is
consistent with earlier estimates due to the lower ISM density
estimate.
The enhanced elemental abundances and high temperature
of the bright emission in the remnant’s eastern half unambigu-
ously demonstrate the presence of hot, metal-rich ejecta. By
contrast, spectral analysis of the emission around the perime-
ter of the ejecta region and throughout the SNR’s western
half (with the possible exception of region 2 — see discus-
sion in § 3.1) reveals cooler material with solar-type com-
position, which supports shocked interstellar/circumstellar
medium (ISM/CSM) for its origin. This well-defined inter-
Chandra Observation of G350.1−0.3 7
Table 3
Spectral Fits to XMMU J172054.5−372652
Parameter BB PL BB+PL BB+BB
NH(1022 cm−2) 3.4+0.2
−0.14.4±0.1 4.1+1.4
−0.64.0+0.4
−0.3
kT1(keV) 0.50 ±0.01 ··· 0.42+0.05
−0.06 0.41±0.05
R1(km) 1.59±0.02 ··· 2.59±0.06 2.94±0.03
Γ··· 5.5±0.3 3+2
−5···
kT2(keV) ··· ··· ··· 0.9+0.8
−0.2
R2(km) ··· ··· ··· 0.19 ±0.01
Fluxa(10−13 ergs cm−2s−1) 4.4±0.1 4.7±0.2 4.9±0.2 4.7±0.1
F1/F2b··· ··· 1.5 9.8
χ2
ν/dof 1.02/105 0.94/105 0.86/103 0.86/103
aThe flux is over the energy range 0.5−10.0 keV and has been corrected for fore-
ground absorption
bUnabsorbed flux ratio between the first and second spectral components in the 1.0−
10.0 keV energy range
10−3
10−2
10−1
1 2 3 4 5
−4
−2
0
2
4
(BB+PL)
energy (keV)
counts s -1 keV -1
Figure 6. The X-ray spectrum of CCO XMMU J172054.5−372652 overlaid
with a BB+PL model. The fit parameters are shown in Table 3.
face between the ejecta and shocked ISM regions further sup-
ports G08’s suggestion that G350.1−0.3 is interacting with a
molecular cloud in the east. If this is indeed the case, the ex-
panding shock wave propagating eastward through the dense
gas would form a strong reverse shock and heat the metal-rich
ejecta to X-ray emitting temperatures while the shock front
propagating in the other direction would still be expanding
relatively unimpeded into the diffuse ISM.
Although we cannot conclusively determine whether the
24 µm emission is due to dust or line emission, the simi-
larity between the infrared and X-ray morphologies supports
the shocked dust scenario (Seok et al. 2008). In addition, the
lack of significant emission at 70 µm and in the IRAC band is
consistent with dust emission that peaks between those bands.
The catalog of 24 µm MIPSGAL sources by Mizuno et al.
(2010) lists only a handful of SNRs, but spectroscopic obser-
vations with Spitzer have shown that SNRs interacting with
dense molecular clouds produce a number of shocked molec-
ular and atomic species in the mid-infrared band (Hewitt et
al. 2009). Thus the presence of 24 µm emission coinciding
with the eastern bright X-ray clump in G350.1−0.3 is con-
sistent with a molecular cloud interaction. As can be seen
from Figure 4, the upper filament (region 1b) appears to out-
line precisely the eastern edge of of the region corresponding
to the ejecta-dominated emission. By contrast, the lower fil-
ament (region 1a) is not visible at 24 µm. The Mg and Si
lines in the spectra of these two regions (Figure 2) as well
as the spectral fit results (Table 1) suggest that the ionization
timescale of the region 1a filament is noticeably higher than
that of the region 1b filament. Although we cannot determine
whether the difference is due to a higher density or longer
time since shock, the material in region 1a (in either case)
may have cooled/sputtered more than that in the region 1b fila-
ment. Incidentally, both regions of emission at 24 µm (1c and
4) coincide with the regions of lowest ionization timescale, as
determined from their X-ray spectra. This result seems con-
sistent with the scenario that the 24 µm emission is due to
dust heated by the reverse shock propagating back into the
SNR and the difference in brightness between the various re-
gions could be attributed to dust cooling or sputtering.
If the explosion occurred approximately at the center of
the X-ray emission (RA. 17:21:01.019, DEC. -37:26:49.61),
the derived age puts the projected velocity of XMMU
J172054.5−372652 at 1400 −2600 km s−1, assuming it was
formed in the same explosion. The lower limit of this veloc-
ity, although high, falls within the 1100 −1600 km s−1mea-
sured for the CCO in Puppis A (Hui & Becker 2006: Win-
kler & Petre 2007). The upper limit is significantly higher
than that measured for any CCO. Although the spectrum of
XMMU J172054.5−372652 can be fitted with a BB+BB or a
BB+PL model, the former requires a BB component of unusu-
ally high temperature. On the other hand, a BB+PL model is
not uncommon when comparing to the X-ray spectra of other
CCOs. The best-fit BB+PL parameters in Table 3 fall within
the typical range of others (see Table 2 in Pavlov et al. 2004).
Indeed, XMMU J172054.5−372652 could be a close cousin
of the CCOs inside SNRs Vela Jr (G266.1-1.2) and Puppis A.
These objects have an age of 1 −3 kyr and a blackbody tem-
perature of ∼0.4 keV. We also note that for the CCO inside
Puppis A, XMM-Newton observations indicate an additional
hard spectral component that can be fitted with either a PL of
Γ= 2.0−2.7 or a hard BB with kT = 0.5−1.1 keV (Becker &
Aschanbach 2002), very similar to our case.
In addition to the prominent difference in flux between the
bright eastern emission and the surrounding regions, there
is an intriguing contrast in brightness between the ambient
material in the far west and the bright regions of shocked
CSM/ISM (regions 3 and 4 in Figure 1a). We estimate the
brightness ratio of these regions to be ∼3, implying a density
ratio of ∼1.7. Thus, if more of the remnant lies unseen fur-
ther to the west, it would need to be only slightly less dense
8 Lovchinsky et al.
to emit below the background. If this is indeed the case, it
would change all the derived SNR parameters and imply that
the SNR may have already entered the ST phase. The larger
radius would increase the SNR’s age, lower the shock velocity
and lower the proton temperature. In addition, it would im-
ply that XMMU J172054.5−372652 may be closer to the ex-
plosion center than evident from the visible emission, which
would place its speed at a lower, more typical value. On the
other hand, if G350.1−0.3 is indeed interacting with a molec-
ular cloud along its eastern edge, the apparent center of the
emission is likely a poor estimate of the explosion location.
In this case, the supernova center would be closer to the east-
ern edge, which would tend to increase the velocity of the
CCO.
5. CONCLUSION
In this paper, we have investigated the dynamical prop-
erties of G350.1−0.3 and its candidate neutron star XMMU
J172054.5−372652. By doing spectral modeling on the high
resolution X-ray data, we have found that the region of bright-
est emission is dominated by hot, metal-rich ejecta while the
diffuse material throughout the surrounding regions is mostly
cooler swept-up CSM/ISM with solar abundances. These re-
sults further support the conclusion of Gaensler et al. (2008)
that G350.1−0.3 is interacting with a dense molecular cloud
in the east. The X-ray imaging has resolved new morpho-
logical features and revealed that the SNR is far more ex-
tended than apparent from the ejecta region alone. We used
the numerical model of Truelove and McKee (1999) to show
that G350.1−0.3 is likely in the free expansion phase and de-
rive values for its age, shock velocity and proton temperature.
We examined the relationship between the inferred shock ve-
locity and the derived electron/proton temperature ratio for
G350.1−0.3 and found it to be entirely consistent with that
of other SNRs. The derived age puts the speed of XMMU
J172054.5−372652 at an unusually high value although its
spectral characteristics are found to be consistent with CCOs
in other remnants. We presented an archival 24 µm image
of G350.1−0.3 and found that the infrared and X-ray mor-
phologies are well-correlated, offering additional support for
a molecular cloud interaction.
The authors would like to thank Tea Temim and Daniel Cas-
tro for their input and several helpful discussions. I.L. ac-
knowledges support from Chandra grant GO9-0059X. P.O.S.
acknowledges partial support from NASA contract NAS8-
03060. J.D.G is supported by an NSF Astronomy and Astro-
physics Postdoctoral Fellowship under award AST-0702957.
B.M.G. acknowledges the support of the Australian Research
Council through grant FF0561298.
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