# Locating the CSM Emission within the Type Ia Supernova Remnant N103B

Benson T. Guest Department of Astronomy, University of Maryland, College Park,
MD 20742, USA  NASA Goddard Spaceflight Center, Greenbelt, MD 20771, USA
Center for Research and Exploration in Space Science and Technology,
NASA/GSFC, Greenbelt, MD 20771, USA  William P. Blair The William H. Miller
III Department of Physics and Astronomy, Johns Hopkins University, 3400 N.
Charles Street, Baltimore, MD 21218, USA  Kazimierz J. Borkowski North
Carolina State University, Raleigh, NC 27607, USA  Parviz Ghavamian Department
of Physics, Astronomy and Geosciences, Towson University, Towson, MD 21252,
USA  Sean P. Hendrick  Millersville University, Millersville, PA 17551, USA
Knox S. Long Space Telescope Science Institute, 3700 San Martin Drive,
Baltimore MD 21218, USA  Eureka Scientific, Inc. 2452 Delmer Street, Suite 1,
Oakland, CA 94602-3017  Robert Petre NASA Goddard Spaceflight Center,
Greenbelt, MD 20771, USA  John C. Raymond Harvard-Smithsonian Center for
Astrophysics, Cambridge, MA 02138, USA  Armin Rest Space Telescope Science
Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA  Ravi Sankrit Space
Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA
Ivo R. Seitenzahl University of New South Wales, Australian Defence Force
Academy, Canberra, ACT 2600, Australia  Brian J. Williams NASA Goddard
Spaceflight Center, Greenbelt, MD 20771, USA

###### Abstract

We present results from deep Chandra observations of the young Type Ia
supernova remnant (SNR) 0509-68.7, also known as N103B, located in the Large
Magellanic cloud (LMC). The remnant displays an asymmetry in brightness, with
the western hemisphere appearing significantly brighter than the eastern half.
Previous multi-wavelength observations have attributed the difference to a
density gradient and suggested circumstellar material origins, drawing
similarities to Kepler’s SNR. We apply a clustering technique combined with
traditional imaging analysis to spatially locate various emission components
within the remnant. We find that O and Mg emission is strongest along the
blast wave, and coincides with Spitzer observations of dust emission and
optical emission from the non-radiative shocks. The abundances of O and Mg in
these regions are enhanced relative to the average LMC abundances and appear
as a distinct spatial distribution compared to the ejecta products, supporting
the circumstellar medium (CSM) interpretation. We also find that the spatial
distribution of Cr is identical to that of Fe in the interior of the remnant,
and does not coincide at all with the O and Mg emission.

##  1 Introduction

Type Ia supernovae are the result of a thermonuclear detonation of a white
dwarf star at the Chandrasekhar mass. The path to explosion is likely to have
multiple origins. The simple picture is either a single degenerate system
where a white dwarf accumulates mass from a non-degenerate companion star, or
the merger of two white dwarfs in a double degenerate system. Each of these
paths has numerous theoretical sub-types and possibilities (e.g Maoz et al.
(2014)). Proving which origin a particular remnant has, and determining the
dominant Type Ia channel remains an active area of investigation. One approach
is to look for evidence of interaction with a circumstellar medium (CSM)
remaining from mass outflow from the binary system prior to the explosion.
This might be expected from a single degenerate progenitor system, but not
from a double degenerate. Another method is through analysis of ejecta
abundances and structure.

Here we present results from a deep Chandra observation of the supernova
remnant (SNR) 0509–68.7, often labelled as N103B, located in the Large
Magellanic Cloud (LMC). Spitzer observations (Williams et al. (2014)) suggest
that the remnant is interacting with a dense CSM (n${}_{0}\sim 10$ cm-3),
drawing similarities to Kepler’s SNR which shows comparable densities (Blair
et al. (2007); Williams et al. (2012)). Li et al. (2021) used Multi-Unit
Spectroscopic Explorer (MUSE) observations obtained with the Very Large
Telescope (VLT) to study the dense knots of CSM finding densities $\geq
10,000$ cm-3. Currently, these are the only two known Type Ia remnants which
show dense CSM interaction hundreds of years after the explosion.

N103B spans 30′′ in diameter, which corresponds to 7.2 pc at the LMC distance
of 50 kpc. Hughes et al. (1995) used X-ray observations with ASCA to identify
N103B as the remnant of a Type Ia supernova, overturning a previous assumption
of a core-collapse progenitor by Chu & Kennicutt (1988) based on the proximity
to an H II region and the OB association NGC 1850. The Type Ia identifier has
been supported by numerous further studies (Lewis et al. (2003); Lopez et al.
(2011); Yang et al. (2013); Yamaguchi et al. (2014)). The star formation
history surrounding N103B was studied by Badenes et al. (2009), finding
vigorous star formation in the recent past suggesting the progenitor or its
companion may have been relatively massive. A light echo from the SN was found
by Rest et al. (2005), from which they derived an age of $\sim 860$ years.
Further support for this young age is provided by expansion measurements from
Chandra observations separated by 17 years (Williams et al., 2018), where the
overall average shock velocity is found to be in excess of 4000 km s-1, and
measurements of balmer dominated shocks with average velocity of $2070\pm 60$
km s-1 by Ghavamian et al. (2017). The Balmer dominated filaments are in
locations with significant preshock neutral hydrogen. The difference between
the Balmer line velocity and measured X-ray proper motion velocity is
consistent with measurements of other SNRs e.g. Tycho with Balmer line
velocities of $\sim 2000-3000$km s-1 and X-ray proper motion velocities of
$\sim 4000-6000$km s-1 (Ghavamian et al. (2001); Williams et al. (2017)).

The X-ray spectrum of N103B, shown in Figure 1, shows strong lines from the
“standard” elements typically seen in Type Ia SNR spectra (Si, S, Ar, Ca, Fe).
Lines from O and Mg are seen as well. Whether the O and Mg results primarily
from the CSM or from unburnt ejecta remains an open question. Burkey et al.
(2013) developed an X-ray analysis technique utilizing a data clustering
algorithm to identify the location of components characteristic of CSM and
ejecta within a remnant. They applied their technique to Kepler’s SNR,
locating within the remnant the areas with pronounced O and Mg emission and
showing that they are spatially coincident with CSM emission, rather than
ejecta. With such obvious similarities between Kepler and N103B Williams et
al. (2014), an obvious question to ask is whether the O and Mg seen in N103B
originates from the CSM or the ejecta.

We apply a modified version of this technique to the N103B observations,
combined with traditional imaging and spectroscopic analysis. We show that the
most significant O and Mg emission appears to follow a distinct spatially
separated distribution from the intermediate mass ejecta dominated emission.
Additionally, we show that the distribution of the O and Mg emission
morphologically correlates with emission from the CSM observed in other
wavelengths.

##  2 Observations

N103B was observed early in the Chandra mission for $\sim 40$ ks (PI: G.
Garmire) in 1999. On February 3, 2003 a $\sim 30$ ks observation (PI: S.
Portegies Zwart) had N103B fall within the field of view; however, the target
was a LMC star cluster, and N103B is far off axis where Chandra’s spatial
resolution is substantially lower. As a result, these data are not used for
spectral analysis in this work. A 400 ks program (PI: B. Williams) was split
into 12 segments and completed between March 20 and June 1 2017. The details
are given in Table 1 with all observations using the ACIS-S array, and the
1999 and 2017 observations positioning N103B on Chandra’s optical axis
pointing of the S3 chip. The data were processed using the CIAO version 4.11
chandra_repro script and CalDB version 4.8.3. For spectral modeling we fit the
observations simultaneously, finding strong residuals from the 1999 and 2003
observations compared to the 2017 observations in the soft $\lesssim 1$ keV
range. There appears to be a systematic difference between the epochs, where
the best fit models over predict the flux in the early observations. A future
study may investigate this difference to determine whether there is conclusive
evidence for intrinsic brightening of the source, but we consider the most
likely source of this discrepancy is the contaminant buildup on the ACIS
detectors that primarily affects the softer energies (Marshall et al. (2004)).
For the purposes of this paper, we will therefore use only the 2017
observations for spectral fitting which provide the largest sample of self-
consistent data and exclude those from the earlier epochs to avoid
unnecessarily increasing our systematic errors. We retain the early
observations for our imaging analysis.

ObsID | Year | Exp Time (ks)
---|---|---
125* | 1999* | 36.7*
3810* | 2003* | 29.7*
18018 | 2017 | 39.5
18019 | 2017 | 59.3
18020 | 2017 | 27.2
19921 | 2017 | 16.9
19922 | 2017 | 41.4
19923 | 2017 | 58.3
20042 | 2017 | 19.8
20053 | 2017 | 11.2
20058 | 2017 | 43.8
20067 | 2017 | 29.7
20074 | 2017 | 31.2
20085 | 2017 | 14.9
Table 1: Breakdown of the Chandra observations totalling 459.6 ks. *The 1999
and 2003 observations are not used in the spectral analysis.

##  3 Methods & Results

###  3.1 Comparison with Other Young Type Ia SNRs

We extracted a global spectrum from a circle encompassing the $\sim
30^{\prime\prime}$ diameter remnant, with a background taken from a 27-35′′
radius annulus surrounding the remnant. The spectrum is plotted in black in
Figure 1 along with Chandra spectra of Kepler’s SNR and Tycho’s SNR in red and
blue, respectively. The spectra are arbitrarily scaled to be normalized at 2
keV for display purposes.

These three remnants are all the result of Type Ia SNe in the last $\sim 1000$
yr, yet their spectra differ substantially. N103B and Kepler show remarkable
agreement, particularly compared to Tycho. Kepler and N103B are roughly an
order of magnitude brighter at 1 keV than Tycho is, relative to the rest of
the spectrum. This is similar to the situation in the mid-IR, where Spitzer
observations of these remnants also show that Kepler and N103B are much more
luminous than Tycho (Blair et al. (2007); Williams et al. (2013, 2014)),
despite being at a relatively comparable evolutionary state. The simplest
explanation of the major differences in the IR and X-ray spectra from these
remnants is that Kepler’s SNR and N103B are expanding into very dense
material, two orders of magnitude higher than Tycho. The subtle differences in
Ar and Ca line emission at $\sim 3.1$ and $\sim 3.9$ keV respectively, and the
energy centroid of the Fe K$\alpha$ line at $\sim 6.5$ keV are explained by
the age of N103B being roughly twice that of Kepler, leading to differences in
ionization equilibrium.

Figure 1: Integrated spectra of N103B (black), Kepler (red), and Tycho (blue). The N103B and Kepler spectra have been scaled up to align the 2.1-2.2 keV continuum with that of the Tycho spectrum. Band | Energy (keV)
---|---
O Low | 0.6–0.64
O Line | 0.64–0.72
O High | 0.72–0.78
Fe Peak | 0.73–1.2
Mg Low | 1.2–1.29
Mg Line | 1.29–1.4
Mg High | 1.4–1.5
Si Low | 1.6–1.75
Si Line | 1.75–1.95
Clustering Si High | 1.95-2.2
Si Ly$\alpha$ | 1.95–2.1
Si High | 2.1–2.3
S Low | 2.1–2.3
S Line | 2.3–2.59
S High | 2.59–2.75
Ar Low | 2.94–3.03
Ar Line | 3.03–3.25
Ar High | 3.25–3.4
Ca Low | 3.65–3.8
Ca Line | 3.8–4.0
Ca High | 4.0–4.2
Cr Line | 5.45–5.75
Fe Low | 6.0–6.25
Fe Line | 6.25-6.8
Fe High | 6.8–7.05
Table 2: Band Energies used for Clustering and Equivalent width line images.

###  3.2 Narrow Band Images

We constructed a single image by reprojecting the 14 observations to a common
axis and merging the event files using the CIAO script merge_obs. From this,
we extracted narrow band images using the CIAO tool dmcopy with 1′′ binning
for each of the band energies listed in Table 2. In Figure 2 (left) we show
spatial separation of O and Mg emission from the emission from Si. Distinct
pockets of red, green, and blue are visible. The right panel shows the
emission from S, Ar, and Ca where no clear spatial separation is apparent.
This may point to the intermediate mass elements originating from bordering
burning layers leading to similar distributions of ejected material. The
emission from the ejecta elements follows a different distribution from the
strongest O and Mg emission. We note that the O band might contain some
emission from Fe L-shell lines as well. Chandra’s spectral resolution is not
sufficient to disentangle O from Fe L-shell lines, but for the purposes of
this paper, we simply label the region “O.”

Figure 2: Left: Red: 0.5-0.8 keV (O), Green: 1.2-1.4 keV (Mg), Blue: 1.75-1.95
keV (Si), Right: Red: S, Green: Ar, Blue: Ca. (Band energies listed in Table
2.) The distinct areas of different colors in the left image highlight the
different distributions of the emitting material. The image on the right shows
no distinct regions where one band dominates over the other two, highlighting
the uniformity of the ejecta elements.

###  3.3 Equivalent Width Images

We then constructed equivalent width images (EWIs) which highlight line
emission relative to the strength of the continuum using the procedure
outlined in Hwang et al. (2000) and Winkler et al. (2014). From the merged
event file we created narrow band images binned with 1′′ pixels for each of
the energy bands listed in Table 2. We estimate the continuum by linearly
interpolating between the normalized low and high energy bands to calculate a
predicted continuum image. We then subtracted this continuum image from the
normalized line image and the excess was then divided by the predicted
continuum, and smoothed the resulting image with a 2-pixel Gaussian.

Figure 3 shows the resulting EW images. The O and Mg images show enhancements
that appears as a ring tracing the edge of the remnant with minimal
contributions from the center. The Si image shows enhancement filling the
remnant; however, it is most significant in a horseshoe morphology with lesser
enhancement in the south extending to the center. The Si Ly$\alpha$ line is
strong enough to create a similar image and shows stronger enhancement in the
fainter eastern half. S, Ar, and Ca all show most significant enhancement on
the eastern side with pockets filling the remnant. There is a notable hole in
each image which coincides with the pocket of interior enhancement seen in the
Mg image as the inward extension of the contour in the north. The Fe K$\alpha$
image shows prominent emission filling the remnant with the brightest emission
located to the south west of the remnant center.

Figure 3: Equivalent width line images with the contours containing the
highest intensity regions overlaid. See the text for details.

To examine the correlation between each element we extracted spectra from the
contours highlighting the enhanced regions shown in Figure 3 and compared to
the spectra from a circle encompassing the full remnant with the contour
regions excluded (Figures 4 & 5). To highlight the differences between the
spectra rather than the overall brightness, scaling was changed where
necessary in order to align the spectra at the level of the 2.1–2.2 keV
continuum.

Figure 4: Comparison of spectra from the EWI contours (Figure 3) shown in
black and the rest of the remnant with these regions excluded shown in red.
The spectra have been scaled with a multiplicative constant to align the
2.1-2.2 keV continuum in order to highlight the differences with the exception
of the O Contours which have been scaled to align the 1keV peak in order to
show the presence of the O and Mg lines which are weak or missing in the rest
of the remnant. Figure 5: Comparison of the Fe K$\alpha$ rich contour regions
from Figure 3 (black) and the rest of the remnant (red). The left and right
panels highlight differences at low and high energies.

Primary differences occur in the soft $\sim 0.5-1.5$ keV range. The regions
which show enhanced O lines correlate with the enhanced Mg regions such that
each show prominent lines of the other. The heavier intermediate mass elements
- Si through Ca - do not show any obvious correlation with enhancement of
other elements, nor does the Fe L-shell or K$\alpha$ emission.

We attempted to create an EWI for Cr ($\sim 5.6$ keV), but the line emission
is insufficient to create an image with meaningful statistics given the
minimal apparent line enhancement over the dim continuum emission at 5 keV.
However, comparing images filtered by the line energies of Cr and Fe from
Table 2 (Figure 6) we find a remarkable spatial agreement between the two,
with both displaying prominent emission from a bright knot in the south west
interior.

Figure 6: Cr (left) and Fe K$\alpha$ (right) narrow band line images. The
pixel size is 2′′ and the images have been smoothed with a 2 pixel Gaussian.
The color scale is intensity.

###  3.4 Clustering

A method to identify and spatially separate emission from different origins
was developed by Burkey et al. (2013). Their work focused on separating
emission from shocked CSM from that arising from ejecta. We build on this
method and introduce the adaptive binning algorithm Contbin (Sanders (2006),
111https://github.com/jeremysanders/contbin/). This routine takes an input
image and starting from the brightest pixel builds up a region following the
surface brightness until the number of counts it contains meets a specified
value. The result is a set of puzzle-piece like regions which cover an image
while all containing a specified minimum number of counts. The input image we
used was a reprojected and merged collection of all 14 Chandra observations
binned to 0.25′′ pixels and spanning 0.3-10 keV. The Contbin routine was then
run on this image to create regions containing $\sim 7000$ photon counts. This
produced over 200 regions (Figure 7).

Figure 7: Broadband Chandra input image (left) and the Contbin generated
region map (right). The grayscale in the region map indicates only region
number and is related to the surface brightness and the order in which the
regions were created by the binning algorithm. Spectra from the three small
colored regions are displayed in Figure 8.

For each region, we summed the number of photon counts belonging to narrow
band energies (Table 2). The selected bands are O, Mg, Si, S, and the $\sim 1$
keV Fe complex. We normalized the number of counts in each band by the band
energy width, and divided by a normalized associated continuum band to remove
inherent bias due to continuum brightness. These continuum bands are the O
low, Mg high, Si high, and S high bands. The Si and S lines counts were summed
to produce 1 value, and the Fe complex line used the Si and S high bands
summed for its continuum ratio. For each region we calculated a 4-D vector
containing the O, Mg, Fe peak, and Si plus S ratios. Further processing was
done by normalizing each vector component relative to the mean. We applied the
kmeans clustering algorithm (Lloyd (1982)) to the set of vectors for a set
number of clusters from 2-10. The number of clusters represents the number of
components we wish to identify. The algorithm attempts to minimize the sum of
the the 4D distance from the cluster center to each region associated with
that cluster.

We extracted spectra from each region with background taken from an annulus
surrounding the remnant. These individual regions display varying
characteristics (Figure 8). We merged the spectra from regions identified as
belonging to a common cluster using the CIAO routine combine_spectra, and
examined by eye the combined spectra (Figure 9) to search for the presence of
O and Mg line features. We value a cluster at 1 if both O and Mg features are
seen, 0.75 if 1 is seen and it is unclear whether the other is present, 0.5 if
only one is present, and 0 if neither feature is visible. Results are
displayed in Figure 10.

Figure 8: Sample spectra from 3 individual regions. Blue: bright south west
knot segment (cyan region in Figure 7), Red: eastern arc segment, Green:
Western limb segment. The regions are shown in Figure 7. The specific regions
are arbitrary, however the selected sample show characteristic differences
between spectra of the different clusters such as the low energy excess
visible in the green spectrum and the energy of the peak intensity. Each
spectrum contains over 7000 spectral counts and the binning shows how the
signal to noise decreases at energies higher than the S band ($>2.8$ keV), yet
allows for comparison between the O, Si, S, Mg, and 1 keV bands. Figure 9:
Comparing the merged spectra from 2 different clusters derived from the 10
cluster trial. Note the absence of O and Mg features in the red spectrum.
Figure 10: Visualization of the presence of O and Mg features in the merged
spectra from each of the clustering runs. For display purposes a spectrum is
valued 1 when both O and Mg features are clearly visible, 0.75 when one of the
two is visible and the other is weakly visible, 0.5 when only one of the two
features is visible, and 0 when neither are present. The 2 cluster tile is
colored differently since both clusters show similar O and Mg features. When
larger number of clusters are used, allowing for more subtle differences in
the spectra to become important the O and Mg features are found to be most
visible on the outer edge of the remnant.

In agreement with the EWIs, the regions bordering the edge of the remnant show
prominent O and Mg line features while the interior regions, particularly for
cluster numbers greater than 6, show either Mg and no O or neither of the two
features. An additional result was the identification of clusters which
contained strong line wings belonging to additional Si and S transitions, such
as the Si Ly$\alpha$ line at 2.01 keV. See Figure 11 for an example.

Figure 11: Sample of the merged spectra from 2 of the clusters derived from
the 9 cluster run. There is a notable difference in the strength of the Si
Ly$\alpha$ line at 2.01 keV.

We find that for larger numbers of clusters, spectra continue to separate into
groups with similar characteristics. With 7 or more clusters, the O and Mg
features are found in spectra only derived from regions bordering the edge of
the remnant, with the most prominent features found on the western side. The
strongest Ly$\alpha$ lines of Si and S are found from a band which runs from
the bottom of the remnant towards the center.

We investigated these regions in more detail by fitting each of the merged 10
clusters spectra with a VNEI model over the limited energy range $1.7-2.8$
keV, a range that covers the Si and S lines. Williams et al. (2017) showed
that with strong emission lines from Si and S, this range is sufficient to
constrain the ionization state of the plasma. The ionization timescale is
mapped in Figure 12. The ionization timescale correlates inversely with
brightness, with the bright western half of the remnant displaying a smaller
timescale than the fainter eastern half. This is counter intuitive since the
ionization timescale should correlate with electron density and brightness. We
investigated the constraints on the ionization fits with the error contour
plots in Figure 13. The ionization timescale is not tightly constrained for
the fainter central region, yet it favours a larger timescale and smaller
temperature. The large difference in brightness between the eastern half of
the remnant and the region we identify with the largest ionization timescale
likely plays a role. The excess continuum emission unaccounted for in our
1-component model may effectively weaken the appearance of the Ly$\alpha$ line
leading to a systematically lower timescale in the brighter regions. Despite
these uncertainties, this large ionization timescale region is spatially
coincident with a similar region identified as an inner ring by the component
analysis of Yamaguchi et al. (2021) supporting the presence of substructure
within the remnant.

Figure 12: Map of the ionization timescale in units of $10^{10}$s cm-3. The
individual input spectra correspond to the groups identified as similar by the
kmeans algorithm for the 10 clusters run. See the text for a description of
the model. Figure 13: Error contours for the two clusters which found the
smallest (left) and largest (right) ionization timescale from Figure 12.

##  4 Spectral modeling

The EWIs presented in Figure 3 highlight regions where the line emission is
the strongest relative to the continuum. The spectra extracted from contours
surrounding the brightest line emitting regions shown in Figures 4-5 shows
that the EWIs represent real features in the underlying spectra. We cannot
immediately attribute areas of line brightness revealed by the EWIs to locally
enhanced abundances in these regions. It is possible that they may result from
the effect of ionization timescale and plasma temperature in enhancing the
emission. To examine the underlying abundances in more detail, we fit the
spectra with an absorbed multi-component non-equilibrium ionization plasma
model VNEI (Borkowski et al. (2001)). The absorption from the Milky Way is
modeled by a Tbabs model with Wilms abundances (Wilms et al. (2000)) fixed to
$6\times 10^{20}\ \text{cm}^{-2}$ (Dickey & Lockman (1990)) and a free to vary
LMC component modeled with Tbvarabs with the LMC average abundances from
Dopita et al. (2019). Fixing one vnei component with the average LMC
abundances, a second component composed of only the intermediate mass elements
Si, S, Ar, and Ca present and allowed to vary, and a 3rd Fe only component. We
process the observations individually, creating response files and background
for each.

Comparing the best fit models from the EWI contour spectra for Mg, we find
enhanced O and Mg from the outer ring contours compared to the interior. The
best-fit models from the Mg enhanced contours and the remnant with these
regions excluded are listed in Table 3, and the spectra are shown in Figure
14.

Component | Parameter | Contours | Excluded
---|---|---|---
Galactic | nH ($10^{22}$ cm-2) | 0.06 (Frozen) | 0.06 (Frozen)
LMC | nH($10^{22}$ cm-2) | 0.15 (0.147 - 0.153) | 0.265 (0.26 - 0.27)
CSM | kT (keV) | 0.9 (0.89 - 0.91) | 0.834 (0.830 - 0.835)
| O | 1.9 (1.87 - 1.95) | 0.63 (0.61 - 0.66)
| Mg | 0.69 (0.68 - 0.71) | 0.42 (0.41 - 0.43)
| Tau ($s$ cm-3) | 1.56E+11 (1.54E11 - 1.58E11) | 1.84E+11 (1.81e11 - 1.88e11)
| Norm (cm-5) | 5.54E-03 (5.51e-3 - 5.55e-3) | 9.72E-03 (9.70e-3 - 9.74e-3)
Ejecta | kT (keV) | 3.32 (3.30 - 3.39) | 4.6 (4.5 - 5.0)
| Si | 2.58 (2.55 - 2.61) | 2.07 (2.01 - 2.15)
| S | 2.43 (2.39 - 2.48) | 1.99 (1.95 - 2.05)
| Ar | 2.74 (2.62 - 2.89) | 2.05 (1.95 - 2.16)
| Ca | 4.5 (4.29 - 4.76) | 3.43 (3.3 - 3.62)
| Tau ($s$ cm-3) | 5.32E+10 (5.24e10 - 5.40e10) | 5.01E+10 (4.94e10 - 5.07e10)
| redshift | -3.80E-03 (-3.395e-3 - –3.8e-3) | -2.72E-03 (-2.77e-3 - -2.68e-3)
| Norm (cm-5) | 2.40E-03 (2.39e-3 - 2.42e-3) | 2.78E-03 (2.72e-3 - 2.83e-3)
Fe | kT (keV) | 7.4 (6.8 - 8.1) | 10.9 (10.7 - 11.4)
| Tau ($s$ cm-3) | 7.03E+10 (6.86e10 - 7.20e10) | 5.66E+10 (5.59e10 - 5.69e10)
| redshift (tied) | -3.80E-03 | -2.72E-03
| Norm (cm-5) | 1.26E-03 (1.23e-3 - 1.29e-3) | 2.10E-03 (2.06e-3 - 2.14e-3)
| $\chi_{\nu}(\nu)$ | 1.787 (2718) | 2.346 (2874)
Table 3: Fits to the spectra extracted from the Mg EWI contours and the entire
remnant with the contours excluded. The spectra are plotted in Figure 14. See
the text for details.  Figure 14: Spectral models from Table 3. The black
spectrum is from the Mg contours and the red spectrum is from the rest of the
remnant. The spectra have been merged for display purposes and the red
spectrum has scaled down by a factor of 3 to avoid overlap. The model derived
from simultaneously fitting the individual observations is overlaid on the
merged spectrum. We have highlighted in grey the O and Mg bands which show
prominence in the contour spectrum.

##  5 Multiwavelength Observations

The multiwavelength picture of N103B reveals a complex structure (Figure 15).
At X-ray energies, the remnant shows a circular morphology with the western
half appearing significantly brighter than the east. At infrared wavelengths,
only the western half is visible, with the deconvolved Spitzer 24 $\mu$m image
revealing a limb and knot structure (Williams et al. (2014)). This same
structure is seen at high spatial resolution in narrow band Hubble images from
Li et al. (2017) and Blair et al. (2020), where it is clear the bright
emission knots, visible in F502N [O III], F673N [S II], F164N [Fe II], arise
from radiative shocks being driven into density enhancements. The H$\alpha$
(F657N) emission shows both the bright clumpy structure and a series of faint,
smooth filaments which align with the Spitzer outer shell structure and arise
from non-radiative shocks along the primary shock front. The presence of these
faint non-radiative shocks is an indicator that the shock is encountering at
least partially neutral gas at these positions. This argument is supported by
Ghavamian et al. (2017) using the Wide Field Integral Spectrograph (WiFeS) on
the 2.3 m telescope at the Siding Spring Observatory in Australia to identify
an intermediate-width component in the H$\alpha$ spectrum with a width of 145
km s$-1$. The overall picture is consistent with the Spitzer conclusion that
denser CSM or ISM is being encountered by the shock on the western side of the
SNR. This is the side of the SNR which is brightest in X-rays as well.

There is a clear offset between the distribution of O and Mg in the X-ray band
which shows prominent emission tracing the blast wave, and the ejecta which
fill the interior. The Mg EWI coincides with the deconvolved Spitzer image
(Figure 16). Figure 17 plots [Fe XIV] 5303 Å emission from MUSE (Seitenzahl et
al. (2019)) in red outlining the location of the reverse shock in the remnant.
The narrowband O line as seen by Chandra is shown in blue, and the O EWI in
green. The narrowband O line shows structure in the eastern half of the
remnant overlapping the Fe reverse shock emission. However, this O line energy
band may contain Fe L-shell emission lines.

Figure 15: Multiwavelength images with corresponding Chandra contours. 1:
Chandra broadband image, 2: Spitzer 24$\mu$m deconvolved (Williams et al.
(2014)), 3: HST H$\alpha$, 4: HST [O III]. The red contour is the Chandra
broadband size, the green contours are the corresponding X-ray line contours.
The H$\alpha$ green contour is the 2.1–2.3 keV continuum emission. Figure 16:
Contours from the Mg EWI overlaid on the deconvolved Spitzer 24 $\mu$m image.
Figure 17: Image of N103B showing the reverse shock as traced by MUSE [Fe XIV]
5303 line emission (Seitenzahl et al. (2019)) in red, and the Chandra O EWI
and narrowband line image in green and blue respectively.

##  6 Discussion

The integrated X-ray spectrum of N103B shows prominent lines from Si, S, Ar,
Ca, and Fe. In addition, there are weaker lines from O and Mg. We have used
spatial analysis using EWIs and the clustering technique described above to
show that the O and Mg lines are found predominantly from a shell bordering
the edge of the remnant while the intermediate mass elements fill the
interior.

Despite the overwhelming difference in brightness of the remnant from west to
east, the distribution of ejecta products does not reflect any obvious east-
west asymmetries in the EWIs. Using multiplicative scaling to align spectra by
their continuum emission brightness we have found strong agreement in the
ejecta line strengths across the remnant. The Si Ly$\alpha$ EWI appears to
favor the eastern half of the remnant. However, the density difference
favoring the bright western half likely lessens the appearance of this line in
the west by increasing the continuum level. The clustering analysis reveals a
distinct separation in the location of the strongest O and Mg emission
compared to the ejecta products in agreement with EWIs. We also find structure
in the ionization timescale distribution. We find stronger Ly$\alpha$ lines
characteristic of a larger ionization timescale from a north-south band in the
eastern center of the remnant.

This is in agreement with the results of Yamaguchi et al. (2021) who used a
component analysis technique to reveal similar structure. Where our results
differ is our abundance measurement of the O and Mg rich emission. Allowing
these parameters to vary, we find enhanced O and Mg from regions spatially
coincident with the Spitzer emission (Williams et al. (2014)), and the outer
blast wave of the remnant. The uncertainties on the abundances are likely much
larger than the values calculated from the data statistics in XSpec, and the
spectrum shows residuals at low energies suggesting that these abundances
should be taken rather as upper limits. However, the presence and coincidence
of the strongest O and Mg line emission with the dust emission as seen in the
24 $\mu$m Spitzer image, and the clear difference in spatial distribution
compared to the ejecta products drives our argument that the remnant is
encountering CSM, likely originating from mass loss from the progenitor
system.

The absence of N lines observable with Chandra means we are unable to contrast
the N to O ratios from the bright radiative knot studied in Blair et al.
(2020). Future high spectral resolution X-ray observations with XRISM and
ATHENA will clarify the extent to which Fe L-shell lines are present in the
“O” band, and will allow a more accurate determination of the elemental
abundances. Finally, IR observations with JWST at a spatial resolution
comparable to that of Chandra will allow for a much more detailed
morphological comparison between the emission from dust and gas in the
remnant.

##  7 Conclusions

Using both EWIs and a clustering technique we have found spatial differences
in the distribution of O and Mg and the intermediate mass ejecta elements. We
have confirmed that these regions are spectroscopically different as well. The
O and Mg EWI and clustering images are morphologically similar to the Spitzer
mid-IR emission from warm dust and the H$\alpha$ emission from non-radiative
shocks, both of which arise from the interaction of the forward shock with the
dense CSM, providing some evidence that the O and Mg seen in the X-rays also
arise from a CSM origin.

The ejecta element EWIs show no preference for the bright western limb in
terms of line enhancement, suggesting the difference in appearance is due to
the enhanced density into which the remnant is expanding. The exception is the
emission in the Cr and Fe K$\alpha$ lines, which are primarily concentrated in
a knot of emission in the southwest. The O and Mg EWIs and spectra extracted
from the most prominent regions contained therein present evidence of enhanced
O and Mg relative to the average LMC abundances, suggesting the remnant is
sweeping into CSM originating from mass loss from the pre-supernova system.

Support for this work was provided by NASA through Chandra General Observer
Program Grant SAO G06-17064. B.G acknowledges the material is based upon work
supported by NASA under award number 80GSFC21M0002. WPB acknowledges partial
support from the JHU Center for Astrophysical Sciences during the time of this
work. We thank the referee for their careful reading of the paper which
improved its quality and clarity.

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