# Supernova Ejecta with Crystalline Silicate Dust in the Supernova Remnant MSH
15$-$52

Hyun-Jeong Kim Department of Earth Science Education, Kongju National
University
56 Gongjudaehak-ro, Gongju-si, Chungcheongnam-do 32588, Republic of Korea
Korea Astronomy and Space Science Institute
776 Daedeok-daero, Yuseong-gu, Daejeon 34055, Republic of Korea  Bon-Chul Koo
Department of Physics and Astronomy, Seoul National University
1 Gwanak-ro, Gwanak-gu, Seoul 08826, Republic of Korea  Takashi Onaka
Department of Astronomy, Graduate School of Science, The University of Tokyo
7-3-1 Hongo, Bunkyo-ku, Tokyo 113-0033, Japan

###### Abstract

IRAS 15099-5856 in the young supernova remnant (SNR) MSH 15$-$52 is the first
and only SNR-associated object with crystalline silicate dust detected so far,
although its nature and the origin of the crystalline silicate are still
unclear. In this paper, we present high-resolution mid-infrared (MIR) imaging
observations of the bright central compact source IRS1 of IRAS 15099-5856 to
study the spatial distributions of gas and dust and the analysis of its
Spitzer MIR spectrum to explore the origin of IRS1. The MIR images obtained
with the T-ReCS attached on the Gemini South telescope show a complicated,
inhomogeneous morphology of IRS1 with bright clumps and diffuse emission in
[Ne II] 12.81 µm and Qa 18.30 µm, which confirms that IRS1 is an extended
source externally heated by the nearby O star Muzzio 10, a candidate for the
binary companion of the progenitor star. The Spitzer MIR spectrum reveals
several ionic emission lines including a strong [Ne II] 12.81 µm line, but no
hydrogen line is detected. We model the spectrum using the photoionization
code CLOUDY with varying elemental composition. The elemental abundance of
IRS1 derived from the model is close to that of SN ejecta with depleted
hydrogen and enhanced metals, particularly neon, argon, and iron. Our results
imply that IRS1 originates from the SN ejecta and suggest the possibility of
the formation of crystalline silicate in newly-formed SN dust.

Infrared spectroscopy(2285) — Supernova remnants(1667) — Interstellar
medium(847)

††facilities: Gemini:South (T-ReCS), Spitzer (IRS)††software: Cloudy version
C13.03 (Ferland et al., 2013), MPFIT (Markwardt, 2009)

##  1 Introduction

Silicates are the most common dust species in the interstellar medium (ISM) of
galaxies. While silicate dust in the ISM of our Galaxy is indicated to be
mostly amorphous (Kemper et al., 2004; Gordon et al., 2023), a sign of
crystalline silicate has also been suggested (Do-Duy et al., 2020).
Crystalline silicates have so far been detected in evolved stars and young
stars (e.g., Molster et al., 1999a; Malfait et al., 1999). Crystallization of
amorphous silicate grains is suggested to occur in circumstellar disks
(Molster et al., 1999b), and several hypotheses of the formation of
crystalline silicate have been proposed including the radial mixing in the
disk and shock waves (e.g., Maaskant et al., 2015, references therein).
Crystalline silicate has been also detected in ultraluminous infrared galaxies
with the crystalline-to-amorphous silicate mass ratios of $\sim$0.1 (Spoon et
al., 2006). This suggests that supernovae (SNe) may be a source of crystalline
silicates, leading to a model of evolution of crystalline silicates in
galaxies (Kemper et al., 2011). Kemper et al. (2011) propose a model of
silicate dust evolution in the galaxy that crystalline silicate is formed in
SNe and amorphized by cosmic-ray hits, suggesting that the crystalline
silicate features could be a useful measure of the youthfulness of the galaxy.

However, the formation of silicate dust in SNe is uncertain. Past observations
in near-infrared (NIR) to mid-infrared (MIR) have revealed the newly-formed
dust in the ejecta of core-collapse SNe although the detected dust mass is
smaller by more than two orders of magnitude than the mass theoretically
predicted (e.g, Sugerman et al., 2006; Sakon et al., 2009; Szalai et al.,
2011). Most of the SN dust have not shown the 10 µm silicate feature,
suggesting that they are mostly carbonaceous dust, and only a few SNe show
evidence for the formation of silicate dust (Kotak et al., 2009; Shahbandeh et
al., 2023). On the other hand, an appreciable amount of dust has been detected
in SN 1987A and young supernova remnants (SNRs) from the far-infrared (FIR)
observations (e.g., Matsuura et al., 2011; Chawner et al., 2020; Millard et
al., 2021), but the dust composition could not be derived because strong dust
features are only present in MIR. From the MIR spectroscopic observations of
SNRs, the presence of carbonaceous dust is suggested (e.g., Tappe et al.,
2006; Andersen et al., 2011), and some SNRs show an indication of silicate
dust including unusual non-stoichiometric silicate and metal oxides (e.g.,
Arendt et al., 2014; Temim et al., 2017; Rho et al., 2018). But there has thus
far not been any observational evidence for the presence of crystalline
silicate in SNe or SNRs except for one case, MSH 15$-$52 (G320.4-1.2), for
which its association with the SN explosion is not yet clearly understood.

MSH 15$-$52 is a young, core-collapse SNR with a complex morphology composed
of the central pulsar wind nebula and a large radio emission about 40 pc in
size at a distance of 5.2$\pm$1.4 kpc (Gaensler et al., 1999, 2002). From the
large extent of the SNR compared to the young age of 1,700 yrs estimated by
the central pulsar PSR B1509-58 (Seward et al., 1983), it was suggested that
MSH 15$-$52 is the remnant of Type Ib SN (SN Ib) with a relatively small
amount of SN ejecta and that the progenitor of the SNR was in a binary system
with the O star Muzzio 10 (2MASS J15135520-5907516) which is $\sim$20$\arcsec$
apart from the pulsar (Gaensler et al., 1999). In the SNR, close to the
pulsar, a bright MIR source IRAS 15099-5856 was discovered from Infrared
Astronomical Satellite observations (Arendt, 1991). IRAS 15099-5856 is only
seen at wavelengths longer than $\sim$10 µm and shows a complicated morphology
with a bright central compact source, a surrounding halo of $\sim$1$\arcmin$
radius with knots and spurs, and several extended ($\sim$4$\arcmin$), knotty
arc-like filaments (Figure 1; Koo et al., 2011).

Figure 1: Three-color image of IRAS 15099-5856 produced with AKARI S11 (B, 11
µm), L15 (G, 15 µm), and L24 (R, 24 µm) images, which is adopted from Figure 1
of Koo et al. (2011). The cross ($\times$), diamond, and plus (+) symbols
present the peak position of IRS1 at 15 µm, O star Muzzio 10, and the pulsar
PSR B1509-58, respectively.

Koo et al. (2011) investigated the central compact source (IRS1) of IRAS
15099-5856 with the AKARI MIR imaging observations and the Spitzer IRS
spectroscopy. The absence of emission at short ($\lesssim$10 µm) wave bands
and the extended morphology observed in the AKARI images imply that IRS1 is an
extended source likely heated by a nearby O star Muzzio 10 as proposed earlier
(Arendt, 1991). A unique feature of IRAS 15099-5856 revealed by the Spitzer
IRS spectrum is the prominent crystalline silicate dust features (Koo et al.,
2011), which has raised an intriguing question about the origin of IRAS
15099-5856 because it is the first and only detection of crystalline silicate
associated with SNRs so far. Koo et al. (2011) proposed a scenario that IRS1
is the material from the progenitor of the SNR ejected at its final
evolutionary stage based on the Spitzer spectrum which is well explained by
dust models including Mg-rich crystalline silicates and the proximity among
IRS1, Muzzio 10, and the pulsar PSR B1509-58. In this scenario, IRS1 might
have survived the SN blast wave as being shielded by Muzzio 10, the former
binary companion star of the progenitor. However, the nature of IRAS
15099-5856 and its association with Muzzio 10 and/or the central pulsar are
still uncertain. While the proper motion of Muzzio 10 is known as $4.9215\pm
0.0163$ mas yr-1 with position angle (PA) = 244.2$\arcdeg$ (measured from
north to east) from the Gaia Data Release 3 (Gaia Collaboration et al., 2016,
2023), the proper motion of PSR B1509-58 has only been derived with huge
uncertainties. Gaensler et al. (1999) reported one-sigma upper limits on the
pulsar’s proper motion: 39 mas yr-1 in right ascension and 52 mas yr-1 in
declination. Leung (2018) measured the proper motion of $\mu_{\alpha}=2\pm 12$
mas yr-1 and $\mu_{\delta}=-50\pm 24$ mas yr-1. The proper motion of the
pulsar seems to imply that the site of the SN explosion about 1,700 yrs ago
was $>1\arcmin$ apart from Muzzio 10, i.e., no association between the SN
progenitor and Muzzio 10, but more accurate measurements are required to
confirm.

In this paper, we investigate IRS1 as a follow-up of Koo et al. (2011). In
Section 2, we present the high-resolution MIR imaging observations of IRS1 and
examine the spatial morphology of IRS1 in detail. In Section 3, we analyze the
Spitzer spectrum of IRS1 by model calculations to derive the elemental
abundance of gas and investigate dust emission. Particularly, we take account
of geometry and energy balance to draw a more physically realistic figure of
IRS1. In Section 4, we discuss the origin of crystalline silicate in MSH
15$-$52 and dust formation in SN ejecta. In Section 5, we summarize and
conclude our study.

##  2 Mid-Infrared Observations of IRAS 15099-5856 IRS1

###  2.1 T-ReCS Observations and Data Reduction

We observed the central compact source IRS1 of IRAS 15099-5856 using the
Thermal-Region Camera Spectrograph (T-ReCS; Telesco et al., 1998; De Buizer &
Fisher, 2005) attached on the Gemini South telescope (Program ID:
GS-2012A-C-4; PI: Onaka, T.) on 2012 May 11 UT. T-ReCS uses a Raytheon
$320\times 240$ pixel Si:As IBC array, providing a pixel scale of $0\farcs
089$ pixel-1 with a field of view of $28\farcs 8\times 21\farcs 6$. We applied
the standard chop-nod technique in order to remove time-variable sky
background, telescope thermal emission, and the 1/f noise in detector. The
chop throw and angle were 15$\arcsec$ and 20$\arcdeg$, respectively. Images
were obtained with the Si-6
($\lambda_{0}=12.33~{}\micron,\Delta\lambda=1.18~{}\micron$), [Ne II]
($\lambda_{0}=12.81~{}\micron,\Delta\lambda=0.23~{}\micron$), [Ne II]cont
($\lambda_{0}=13.10~{}\micron,\Delta\lambda=0.22~{}\micron$), Qa
($\lambda_{0}=18.30~{}\micron,\Delta\lambda=1.51~{}\micron$), and Qb
($\lambda_{0}=24.56~{}\micron,\Delta\lambda=1.92~{}\micron$) filters, among
which the Si-6 and [Ne II]cont filters were used to determine the continuum
baseline of the [Ne II] image. For flux calibration, we observed standard
stars $\gamma$ Cru and $\omega$ Lup (HD 139127) from Cohen standards (Cohen et
al., 1999) with the same filters. The total exposure time of IRS1 was 300 sec
for the Qa filter and 900 sec for the others. The standard stars were observed
with the exposure time of 30 sec for all filters.

Data were reduced by using the custom IDL software MEFTOOLS version
5.0111MEFTOOLS was developed and provided by James M. De Buizer via
http://www.jim-debuizer.net/research/, but it is no longer available.. During
the image stacking, bad-frames such as the ones affected by instrumental
artifacts have been excluded via visual inspection. We corrected relative
offsets between filters using standard stars observed by the same order as
IRS1, but it was unable to obtain the accurate astrometric solutions because
IRS1 is an extended source without nearby stars bright in MIR that can be used
as a reference. Instead, we corrected the absolute astrometry using the O star
Muzzio 10 only detected in the Si-6 image. Although this is a rough
correction, the peak position in the Qa-band ($\lambda_{0}=18.30~{}\micron$)
is coincident with the peak position defined based on the AKARI 15 µm image
within $<$0$\farcs$2\. The seeing estimated from the standard stars is about
$0\farcs 6$. For flux calibration, the standard star $\omega$ Lup was used for
all the filters except Qb that used $\gamma$ Cru. The in-band fluxes of
$\omega$ Lup in the Si-6, [Ne II], [Ne II]cont, and Qa filters are 11.381,
10.569, 10.162, and 5.172 Jy, respectively, at airmass of 1, similar to the
airmass of the IRS1 at the time of observations (1.14–1.4); the in-band flux
of $\gamma$ Cru in the Qb filter is 157.567 Jy at airmass of
1222https://www.gemini.edu/sciops/instruments/michelle/MIR-resources.

Figure 2: T-ReCS images of IRS1 compared with the AKARI S11 (11 µm) image. The
cyan contours on the Qa image are the [Ne II] 12.81 µm contours with the flux
levels of 0.3, 0.55, 0.65, 0.9, 1.2, and 1.5 mJy from the outermost. The green
or black contours on the other images are the Qa 18.30 µm contours with the
flux levels of 0.7, 1.1, 1.5, 2.1, and 2.5 mJy. On the AKARI S11 image, the
cross and diamond symbols present the peak position of IRS1 at 15 µm and O
star Muzzio 10, respectively.

###  2.2 Mid-Infrared Morphology

Figure 2 displays the [Ne II], Qa, [Ne II]cont, Si-6, and Qb images of IRS1
obtained with T-ReCS and the AKARI S11 ($\lambda_{0}=11~{}\micron$) image with
the contours of Qa (green, black) and [Ne II] (cyan) overlaid. The T-ReCS
images were smoothed by a Gaussian kernel with a standard deviation of
1$\arcsec$. In the AKARI image with low spatial resolution, IRS1 is
elliptically extended along east-west direction with PA of 110$\arcdeg$. It is
also extended in the T-ReCS images but shows an irregular morphology with sub
structures. In the Qa image, IRS1 consists of two parts in the east and west.
The eastern part is extended along northwest-southeast (PA = 144$\arcdeg$)
direction and composed of three bright clumps although the brightest peaks are
not well defined. The western part, on the other hand, is extended along
northeast-southwest (PA = 50$\arcdeg$) direction and composed of a bright
compact knot and diffuse emission. The size of the bright knot obtained by
Gaussian fitting is about $2\farcs 7\times 1\arcsec$ in FWHM. Owing to this
bright knot, the surface brightness of the western part is comparable to the
brightness of the eastern part (see below) although the western part is about
half the size of the eastern part. The [Ne II] image is overall similar to the
Qa image, but the detailed structure is different. The eastern part in [Ne II]
is extended as large as the eastern part of the Qa image, but two bright knots
are distinctively shown with the size about $1\farcs 4\times 2\farcs 7$ and
$1\farcs 3\times 1\farcs 7$ in FWHM, both of which are larger than the seeing
($\sim 0\farcs 6$) measured from the standard stars. The western part in [Ne
II] is very faint and much smaller than the western part in the Qa image.

A remarkable feature is that the distributions of Qa and [Ne II] emission are
not consistent with each other. Although the whole extent is similar, the
bright peaks have offsets between the two images as shown by the Qa contours
on the [Ne II] image and the [Ne II] contours on the Qa image in Figure 2.
This discrepancy does not likely come from the inaccurate astrometry. The
relative distances between the peaks are also different in the two images.
Therefore, the Qa and [Ne II] images obtained with T-ReCS indicate not only a
complex, inhomogeneous morphology of IRS1 itself but also different
distributions of the gas and dust in IRS1.

The other T-ReCS images besides Qa and [Ne II] do not show any particular
emission. The [Ne II]cont image shows no emission, implying no continuum
emission in the [Ne II] image. There is weak continuum around 13 µm in the
Spitzer IRS spectrum of IRS1 (see Figure 3), but it may be too weak to be
detected in the T-ReCS image. The Si-6 image shows the emission almost
identical to [Ne II]. This implies that there is no other line except [Ne II].
While IRS1 is bright at wavelengths longer than 15 µm with the spectral energy
distribution (SED) peaking at around 30 µm (Figure 3 of Koo et al., 2011), the
Qb image at 24.56 µm in Figure 2 does not detect significant emission because
of lower sensitivity of the Qb filter. In the Qb image, some faint emission
features are shown inside the Qa contours, but they are well below three-sigma
(3$\sigma$) where $\sigma$ ($\simeq 1.0\times 10^{-14}$ erg s-1 cm-2 for Qb)
is an rms noise.

Previously, Koo et al. (2011) suggested that IRS1 is externally being heated
by Muzzio 10, which is 13$\farcs$7 away from IRS1 to the south, based on the
non-detection of an embedded point source in optical/NIR and the temperature
of Muzzio 10 that is appropriate to produce the observed [Ne II] line
luminosity obtained from the Spitzer spectrum (see Section 3 as well). The
T-ReCS images also do not show any signature of a point source embedded in
IRS1, confirming the previous prediction. A star might be deeply embedded in
the bright [Ne II] knots, but non-detection of any stellar source in [Ne
II]cont or Si-6 rules out this possibility.

Table 1: T-ReCS Qa and [Ne II] Flux of IRS1

Filter | Region | Flux | Area | Brightness
---|---|---|---|---
|  | (Jy) | ($\times 10^{-12}$ erg s-1 cm-2) | (arcsec2) | ($\times 10^{-13}$ erg s-1 cm-2 arcsec-2)
| whole | 6.85 | 92.6 | 91.17 | 10.2
Qa | east | 3.84 | 51.9 | 29.07 | 17.9
| west | 1.64 | 22.2 | 13.51 | 16.4
| whole | 2.28 | 9.58 | 71.54 | 1.34
| east | 1.80 | 7.58 | 32.57 | 2.33
[Ne II] | west | 0.15 | 0.63 | 4.47 | 1.42
| east knot | 0.21 | 0.87 | 1.60 | 5.44
| west knot | 0.34 | 1.45 | 2.62 | 5.53

Note. — The uncertainty in flux measurements is $\lesssim$20%.

###  2.3 Flux Estimation

We measured the flux of IRS1 from the Qa and [Ne II] images. Applying aperture
photometry, we estimated the flux of the whole source, eastern and western
parts, and the two bright knots in [Ne II] as listed in Table 1. The source
regions were determined by the $3\sigma$ contours, and the knot regions were
determined by the size of the knots. The uncertainty in flux measurements is
$\lesssim$20%. As described earlier, the Qa flux of the western part is
smaller than that of the eastern part, but the surface brightness of both are
comparable. The [Ne II] flux is concentrated on the bright knots with the
surface brightness twice larger than the other area. For comparison, the Si-6
flux for the same source region as [Ne II] is $1.03\times 10^{-11}$ erg s-1
cm-2, which is a little larger than the [Ne II] flux $9.58\times 10^{-12}$ erg
s-1 cm-2 owing to weak ($\sim 1\times 10^{-12}$ erg s-1 cm-2) continuum at
11.5–13 µm seen in the Spitzer IRS spectrum (see Section 3.1 and Figure 3).

We compare the flux measured from the T-ReCS images with the flux estimated
from the Spitzer IRS spectrum. The Qa flux obtained by using the transmission
curve of the Qa
filter333http://www.gemini.edu/sciops/instruments/trecs/imaging/filters is
10.16 Jy, larger than the Qa flux from the T-ReCS image by a factor of 1.5.
The [Ne II] line flux obtained by Gaussian fitting of the emission line
(Section 3.1; Table 2) is $8.21\times 10^{-12}$ erg s-1 cm-2, about 85% of the
[Ne II] flux from the T-ReCS image. The flux differences between the T-ReCS
and Spitzer observations can be explained by the inhomogeneous morphology of
IRS1 and slit-loss correction of the Spitzer IRS spectrum as well as in part
by sky chopping operation in the T-ReCS observations. The Spitzer IRS spectrum
was obtained with two low-resolution modules: the short-low (SL) module
covering 5.2–14.5 µm and the long-low (LL) module covering 14.0–38.0 µm. The
two slits perpendicularly placed on IRS1 did not cover the same area because
of different slit widths (Figure 1 of Koo et al., 2011), and the SL slit along
north-south direction with a slit width of $3\farcs 7$ only partially covered
IRS1, requiring slit-loss correction. The slit-loss correction factor was
determined by the brightness distribution of the source. IRS1 was assumed as a
Gaussian distribution of $12\arcsec\times 5\arcsec$ in size, which is very
different from the morphology observed in the Qa and [Ne II] images. Since the
LL slit width ($10\farcs 7$) is larger than the size of IRS1 estimated in the
Qa image, the Spitzer spectrum possibly includes extended, diffuse emission as
well that is not detected in the T-ReCS observations, leading to a larger Qa
flux from the spectrum. We note that Koo et al. (2011) derived the slit-loss
correction factor assuming the two-dimensional brightness distribution of IRS1
given by the AKARI images to match the flux between the AKARI images and
Spitzer spectrum, which results in the Qa and [Ne II] flux of 14.3 Jy and
$1.2\times 10^{-11}$ erg s-1 cm-2, respectively.

##  3 CLOUDY Modeling of the Spitzer IRS Spectrum of IRS1

###  3.1 Spectral Characteristics

The Spitzer IRS spectrum of IRS1 presented in Figure 3 was obtained in 2008
October 3 UT (Program ID: 50495; PI: Koo, B.-C.) and examined by Koo et al.
(2011). In the AKARI images from N3 ($\lambda_{0}=3.2~{}\micron$) to L24
($\lambda_{0}=24~{}\micron$), IRS1 is only seen at $>$11 µm. In the Spitzer
IRS spectrum likewise, continuum emission is extremely weak ($<$0.1 Jy) at
$\lesssim$13 µm and steeply increases to $\sim$20 µm. The most remarkable
feature in the spectrum is the strong and relatively narrow peaks at 23, 27,
and 34 µm, which are well explained by crystalline silicate dust. A model of
the IRS spectrum produced by modified-blackbody fitting with dust species
including crystalline olivine (Mg1.9Fe0.1SiO4), metal oxides (FeO, MgO), and
amorphous silicate fairly well reproduces the observed spectrum, providing
total dust mass of $9\times 10^{-3}$ $M_{\sun}$ and dust temperature of
$\sim$55–150 K at a distance of 4 kpc (see Koo et al., 2011, for details).

The spectrum also shows several ionic emission lines including a strong [Ne
II] 12.81 µm line that was partly discussed in Koo et al. (2011). In Figure 3,
the emission lines are not clearly seen because of the prominent dust features
except the [Ne II] line at 12.81 µm. We detected [Ar III] 8.99 µm, [S IV]
10.51 µm, [Ne II] 12.81 µm, [Ne III] 15.56 µm, [S III] 18.71 µm, and [O IV]
25.89/[Fe II] 25.99 µm lines (Figure 3) and measured the line fluxes by the
Gaussian fitting. No hydrogen line has been seen in the spectrum. For the
Gaussian fitting, we used the IDL MPFIT package (Markwardt,
2009)444https://pages.physics.wisc.edu/c̃raigm/idl/fitting.html. The derived
line fluxes are listed in Table 2. The flux errors in the table are from the
Gaussian fitting and do not include the uncertainty of the observed spectrum.
Since the resolving power of the Spitzer IRS LL module555Spitzer Space
Telescope Observer’s Manual version 8.0, Chapter 7.1.6, issued by the Spitzer
Science Center (http://ssc.spitzer.caltech.edu) between 14 and 21.3 µm is
given as $R=2.9524\lambda$, i.e., R $\sim$76 at 25.9 µm, two adjacent lines [O
IV] 25.89 µm and [Fe II] 25.99 µm are not resolved. The [S III] line at 33.48
µm also seems to present in the spectrum, but the line flux was not measured
because the line is severely blended with the strong dust feature at 34 µm.
The FWHM of the emission lines obtained from the Gaussian fitting is from 0.11
to 0.39 µm depending on the wavelength. These line widths are comparable to
the spectral resolving power, implying that the velocity of the ionic lines is
not resolved.

Figure 3: Spitzer IRS spectrum of IRS1 (Koo et al., 2011) with the detected lines marked.  Table 2: Detected Emission Lines and their Fluxes from Spitzer IRS Spectrum of IRS1 Line | Wavelength | Flux
---|---|---
| µm | ($\times 10^{-12}$ erg s-1 cm-2)
[Ar III] | 8.99 | 0.13($\pm$0.02)
[S IV] | 10.51 | 0.14($\pm$0.02)
[Ne II] | 12.81 | 8.21($\pm$0.02)
[Ne III] | 15.56 | 0.97($\pm$0.04)
[S III] | 18.71 | 0.51($\pm$0.07)
[O IV]/[Fe II]aa[O IV] and [Fe II] lines are not resolved at the spectral resolution of the Spitzer IRS LL module. | 25.89/25.99 | 0.93($\pm$0.07)
Figure 4: [Ne III]15.56µm/[Ne II]12.81µm versus [S IV]10.51µm/[S III]18.71µm
line ratio diagram of various astronomical objects (gray cross and open
symbols), Cloudy models of novae (black filled hour glass), and IRAS
15099-5856 IRS1 (red star). The gray lines are the Cloudy model grid produced
with hydrogen density $n({\rm H})$=100 cm-3, heating source temperature from
35,000 K to 50,000 K, and ISM abundance with varying neon abundance. The small
vertical bars in orange, yellow, green, blue, purple, and black indicate neon
abundance from $-4.0301$ to $-1.5301$ with an interval of 0.5 in log scale
relative to hydrogen.

###  3.2 Emission Line Ratios

The IR fine-structure emission lines are frequently used as a diagnostic tool
in the investigations of gaseous nebulae, ionized regions, or obscured clouds.
Particularly, the line ratios of some specific lines have a tight correlation,
providing physical conditions of the region of interest (Dinerstein, 1995;
Dopita & Sutherland, 2003). In Figure 4, we present the line ratio diagram [Ne
III]15.56µm/[Ne II]12.81µm versus [S IV]10.51µm/[S III]18.71µm of various
astronomical objects with the observed line ratios of IRS1 (IRAS 15099). The
line ratios of the objects except novae were obtained from the literature: H
II regions in the Galaxy and Large/Small Magellanic Clouds (LMC/SMC) from
Tables 2 and 5 of Giveon et al. (2002); giant H II regions from Table 2 of
Lebouteiller et al. (2008); LMC/SMC planetary nebulae (PNe) from Table 2 of
Bernard-Salas et al. (2008); a luminous blue variable candidate (cLBV)
G79.29+0.46 from Jiménez-Esteban et al. (2010). The line ratios of novae were
obtained from the model calculations using the one-dimensional plasma
simulation code Cloudy666https://www.nublado.org version C13.03 (Ferland et
al., 2013). Cloudy solves the ionization, chemical, and thermal state of
material exposed to an external radiation field or other heating source, and
predicts observable quantities such as emission and absorption spectra that
can be compared with observations. The nova models were calculated by assuming
a blackbody of $T_{\rm eff}=47,000$ K and $L=6.3\times 10^{36}$ erg s-1
(Schwarz et al., 2007) and by adopting the abundances of the novae V1500 Cygni
(Ferland & Shields, 1978) and V838 Her (Schwarz et al., 2007), which show
enhanced metal abundances (see Table 3). We also overlay a Cloudy model grid
produced with hydrogen density $n({\rm H})$=100 cm-3, heating source
temperature from 35,000 K to 50,000 K, and the ISM abundance with varying neon
abundance from $-4.0301$ (the ISM abundance) to $-1.5301$ in log scale
relative to hydrogen, i.e., $n({\rm Ne})$ from $9.33\times 10^{-3}$ cm-3 to
2.95 cm-3. For the ISM abundance, we adopted the protosolar abundance (Asplund
et al., 2009, see also Table 3). The shapes of the radiation fields (SEDs) of
heating sources were adopted from the pre-calculated stellar atmospheric
models of the Tlusty OB star grids777http://tlusty.oca.eu/Tlusty2002/tlusty-
frames-cloudy.html (Lanz & Hubeny, 2003, 2007) provided along with the Cloudy
code from which we selected the main-sequence star models at solar metallicity
for a given temperature.

It has been known that there is a good correlation between [Ne III]15.56µm/[Ne
II]12.81µm and [S IV]10.51µm/[S III]18.71µm (e.g., Martín-Hernández et al.,
2002). The relation is almost linear, and it suggests that the two line ratios
are almost equally affected by the hardness of the ionizing radiation. In
terms of the stellar $T_{\rm eff}$, the observed line ratios of the H II
regions of low to high ionization structures can be described with $T_{\rm
eff}=$35,000 K to 50,000 K (Figure 4; see also Figure 2 of Martín-Hernández et
al., 2002). Figure 4 shows that the relation is in general consistent with the
theoretical relation expected for H II regions with the ISM abundance,
although the majority of the H II regions with low-ionization structure
appears to be above the theoretical line. In contrast, the nova models are
located well below the theoretical relation for the ISM abundance, which is
likely due to the high abundances of heavy elements. The Cloudy model grid,
for example, demonstrates how the line ratios vary with neon abundance, and it
indicates that the enhanced neon abundance lowers the [Ne III]15.56µm/[Ne
II]12.81µm ratio. The observed line ratios of IRS1 are similar to those of
nova, implying that the elemental composition of IRS1 might be similar to that
of nova. Hence, we adopt the nova abundance as an initial abundance set for
our modeling of the Spitzer IRS spectrum of IRS1 in Section 3.5.

We note that there is an issue about the line ratios of IRS1 derived from the
Spitzer IRS spectrum. As described in Section 2.3, IRS1 was observed with two
IRS modules with different slit widths that covered different parts of IRS1.
Since the two lines of each pair in Figure 4 ([S IV] and [S III]; [Ne III] and
[Ne II]) are from different modules, their uncertainty can be large depending
on the slit-loss correction. We compare the line ratios of IRS1 on the
empirical relation $\rm{log\,([Ne\,{\footnotesize III}]/[Ne\,{\footnotesize
II}])=0.81\times log\,([S\,{\footnotesize IV}]/[Ne\,{\footnotesize
II}])+0.36}$ (Groves et al., 2008) which was derived from the archival spectra
of a wide range of astrophysical objects from nearby H II regions to
ultraluminous infrared galaxies obtained by Spitzer and Infrared Space
Observatory (ISO; Kessler et al., 1996). With the [S IV]/[Ne II] ratio derived
from the same SL module, the [Ne III]/[Ne II] ratio expected by this relation
is 0.09, which is comparable to [Ne III]/[Ne II] $\sim$0.12 derived from the
observed spectrum. Therefore, we assume that the slit-loss correction is
acceptable, although there is still uncertainty from the brightness
distribution between the one we assumed (i.e., 2D Gaussian distribution) and
the real distribution of IRS1.

###  3.3 Model Parameters and Assumptions

Previously, Koo et al. (2011) modeled the Spitzer IRS spectrum as thermal
emission from several independent dust components using the modified
blackbody. Their models well reproduce the observed spectrum, but the derived
dust temperatures show large differences ranging from 55 K to 150 K because
they treated dust components independently. In this study, we model the IRS
spectrum using Cloudy to include physical process and energy balance. We first
set the geometry of model calculation. The observations imply that IRS1 with a
size $9\farcs 6\times 5\farcs 1$ or 0.19 pc $\times$ 0.10 pc at the distance
of 4 kpc (Koo et al., 2011), is externally heated by Muzzio 10 separated by
$13\farcs 7$ or 0.27 pc, in projected distance. The separation between IRS1
and Muzzio 10 ($r_{0}$) in principle should be treated as a free parameter
because it significantly affects the radiation absorbed by IRS1 and total dust
mass, but we fixed it to reduce the number of free parameters based on the
initial models (See Section 3.4). We also fixed the thickness of IRS1 as the
same as the major axis (0.19 pc) of IRS1 on the projected sky. Figure 5 is a
schematic figure of the geometry. Since the heating source (= Muzzio 10) is
outside the cloud (= IRS1), we adopted a covering factor ($=\Omega/4\pi$,
where $\Omega$ is an area of the cloud divided by the distance between the
heating source and cloud.) to take account into a fraction of the radiation
field emitted by the heating source that actually strikes the cloud. In Figure
5, the ‘transmitted’ radiation is the net emission emerging from the shielded
face of the cloud, and it includes both the attenuated continuum radiation of
Muzzio 10 and the diffuse emission emitted from IRS1. The ‘reflected’
radiation is the emission from the illuminated face of the cloud back into the
direction towards the heating source, and it includes both the backscattered
incident radiation of Muzzio 10 and the diffuse emission emitted from IRS1.
The reflected emission corresponds to the emission that we observe.

Figure 5: Geometry of Cloudy model calculations of IRS1. The blue star is the
heating source Muzzio 10. ‘Inc.’ refers to incident continuum from Muzzio 10,
and ‘Trans.’ and ‘Refl.’ refer to reflected and transmitted radiation,
respectively.

The heating source in model calculations was fixed as Muzzio 10. The spectral
type of Muzzio 10 is O4.5III(fp) (M. Bessell 2010, private communication) or
O5n(f)p (Maíz Apellániz et al., 2016). Since the luminosity class is uncertain
for the latter, we adopted the stellar parameters of an O4.5III star (Martins
et al., 2005): $T_{\rm eff}=40,500$ K, ${\rm log}~{}g=3.71$ cm s-2, and log
$L/L_{\sun}$ =5.76 or ${\rm log}~{}Q_{0}=49.52~{}{\rm s}^{-1}$. For the shape
of the radiation field, we used the Tlusty O star model at solar metallicity
with $T_{\rm eff}=40,000$ K and ${\rm log}~{}g=3.75$ cm s-2 (Lanz & Hubeny,
2003).

Dust species were adopted from Koo et al. (2011). We calculated the
absorption/scattering coefficients of each dust from their optical constants
following the Bohren-Huffman Mie scattering (Bohren & Huffman, 1983) and
compiled dust opacity files using the grain code in Cloudy. The optical
constants of dust were adopted from the literature: crystalline olivine
(Mg1.9Fe0.1SiO4) from Fabian et al. (2001); FeO and Mg0.6Fe0.4O from Henning
et al. (1995). For amorphous silicate, we used the opacity provided in Cloudy.
The compiled opacity assumes a spherical dust grain with a size of $0.25$ µm
for FeO and $0.1$ µm for the others. Figure 6 presents the absorption
coefficients ($Q_{\rm abs}$) of each dust species.

Figure 6: Absorption coefficients ($Q_{\rm abs}$) of dust species used in
Cloudy model calculations. The black, red, blue, and green colors represent
crystalline olivine (Mg1.9Fe0.1SiO4), FeO, Mg0.6Fe0.4O, and amorphous
silicate, respectively. The size of a dust grain is 0.1 µm in radius except
FeO with a radius of 0.25 µm.

###  3.4 Constraints from SED Models

Since a large number of free parameters are involved in Cloudy calculations,
we first modeled the SED of IRS1 to constrain some parameters. Besides the
assumed geometry (Figure 5), heating source of an O4.5III star, and dust
species from Koo et al. (2011), low hydrogen density was required to produce
an SED that shows no emission at $\lesssim$11 µm. We also found that
$r_{0}=0.45$ pc would reasonably well fits the observed MIR flux levels, so we
fixed $r_{0}$ as 0.45 pc.

We also produced the models with an internal heating source to verify a
possibility that there is a deeply embedded heating source inside IRS1.
Assuming a central star surrounded by dust cloud at a certain distance
($R_{\rm in}$), we produced the models with heating sources of $T_{\rm
eff}=30,000$ K and 19,000 K corresponding to a B0V and B3V star, respectively.
The other parameters such as hydrogen density, $R_{\rm in}$, or abundance were
adjusted to match the observed MIR flux and a dip around 10 µm. Figure 7
compares the observed and model SEDs reddened by the column density of
$9.2~{}\times~{}10^{21}$ cm-2 with $R_{\rm V}=3.1$ (Koo et al., 2011). In the
models with an embedded heating source, as shown in the figure, continuum
emission in NIR and even in optical wavelengths is definitely expected. While
a dip around 10 µm can be produced and become deeper by larger $R_{\rm in}$,
continuum emission at short wavelengths is always predicted. A heating source
with significantly low temperature can reduce continuum flux but cannot
produce ionic emission lines we observe in the IRS spectrum. For example, the
strong [Ne II] line at 12.81 µm is not produced by a heating source with
$T_{\rm eff}=19,000$ K (green line in Figure 7). From the models with an
internal heating source, we confirm that IRS1 is externally heated by Muzzio
10.

Figure 7: SEDs of IRS1 from the observations and Cloudy models with an
embedded heating source. The blue solid- and red dashed-lines are the models
with the same heating source of $T_{\rm eff}=30,000$ K (B0V) but with
different hydrogen density of ${\rm log}\,n({\rm H})=$ 2.4 and $-4.9$ (or,
$n({\rm H})=250$ and $1.3\times 10^{-5}$ cm-3), respectively. The green line
is the model with the heating source of $T_{\rm eff}=19,000$ K (B3V). The
symbols are the fluxes from various observations as labeled in the legend. The
open symbols with an arrow represent upper limits.

In Figure 7, the blue solid- and red dashed-lines are the models with the same
heating source ($T_{\rm eff}=30,000$ K) but with different hydrogen density of
${\rm log}\,n({\rm H})=2.4$ (or, $n({\rm H})=250$ cm-3) and ${\rm log}\,n({\rm
H})=-4.9$ (or, $n({\rm H})=1.3\times 10^{-5}$ cm-3), respectively. In Cloudy,
densities of gas and dust are defined relative to hydrogen. For these two
models with different hydrogen densities, we adjusted the amount of metals and
grains to result in the same dust masses. The SEDs of the two models are
almost identical in MIR but very different in NIR. While the model with a
moderate hydrogen density (blue) shows a group of emission lines from
hydrogen, the model with negligible hydrogen (red) shows no emission line at
all. This holds true for a case that is heated externally as well (see Figure
10 in Section 3.5). Since we do not observe any hydrogen emission in optical
and/or NIR, we have assumed that hydrogen is depleted in IRS1 and fixed the
hydrogen density to be almost zero (but set the value of $n({\rm H})\sim
10^{-5}$ cm-3 to run the code).

Table 3: Abundances used in the CLOUDY Modeling

Atom | ISMaaFor ISM, the protosolar abundance (Asplund et al., 2009) is adopted. | Nova (V1500 Cyg)bbThe nova abundance stored in Cloudy derived by Ferland & Shields (1978) for V1500 Cygni. | Nova (V838 Her)ccThe nova abundance derived for V838 Her (Schwarz et al., 2007). | IRS1 Model
---|---|---|---|---
H | 1.00E+02 | 3.16E+07 | 3.16E+07 | 1.26E-05
He | 9.55E+00 | 3.09E+06 | 4.47E+06 | 1.23E-06
C | 2.95E-02 | 2.95E+04 | 6.03E+04 | 3.72E-01
N | 7.41E-03 ($-$0.60) | 3.09E+05 (1.02) | 7.41E+04 (0.09) | 3.89E-04 ($-$2.98)
O | 5.37E-02 (0.26) | 5.37E+05 (1.26) | 2.82E+04 ($-$0.33) | 1.07E+00 (0.46)
Ne | 9.33E-03 ($-$0.50) | 6.46E+04 (0.34) | 1.91E+05 (0.50) | 3.24E+01 (1.94)
Mg | 4.37E-03 ($-$0.83) | 1.20E+03 ($-$1.39) | 1.58E+03 ($-$1.58) | 1.51E-02 ($-$1.39)
Si | 3.55E-03 ($-$0.92) | 1.12E+03 ($-$1.42) | 2.04E+03 ($-$1.47) | 1.41E-02 ($-$1.42)
S | 1.45E-03 ($-$1.31) | 5.13E+02 ($-$1.76) | 6.92E+03 ($-$0.94) | 7.24E-02 ($-$0.71)
Cl | 1.86E-05 ($-$3.20) | 5.89E+00 ($-$3.70) | 5.89E+00 ($-$4.01) | 7.41E-05 ($-$3.70)
Ar | 2.75E-04 ($-$2.03) | 1.15E+02 ($-$2.41) | 1.15E+02 ($-$2.72) | 1.15E-01 ($-$0.51)
Fe | 3.47E-03 ($-$0.93) | 1.48E+03 ($-$1.30) | 8.51E+03 ($-$0.85) | 1.17E+01 (1.50)

Note. — The abundance is the absolute number density (cm-3) applied to the
Cloudy models. The numbers in parentheses are the number density of the
element relative to carbon in log scale, which shows the relative abundances
among metals.

###  3.5 Modeling of the Spitzer IRS Spectrum

From the SED models, we fixed the geometry of an externally heated cloud, the
heating source of an O4.5III star, and negligibly-small hydrogen density. With
these parameters, we modeled the Spitzer IRS spectrum to derive the
physical/chemical characteristics of IRS1. In Cloudy, the initial abundance
can be adopted from the stored abundance sets including H II regions, general
ISM, novae, and PNe. Based on the line ratio diagram (Figure 4, Section 3.2),
we adopted the abundance of the nova V1500 Cyg listed in Table 3 as an initial
abundance set and adjusted the amount of metals and grains by changing the
scale factors. The elements not listed in Table 3 were not included in the
calculations. With the hydrogen density ${\rm log}\,n({\rm H})=-4.9$, we
applied a scale factor of 7.5 in log scale to metals (the elements heavier
than helium) and grains in the nova abundance to fit the MIR flux level of the
IRS spectrum.

We first searched for a model that reproduces the observed dust features by
changing the amount of four dust species: crystalline olivine, FeO,
Mg0.6Fe0.4O, and silicate. With the fixed dust abundance, we determined the
gas abundance that explains the observed line flux derived in Section 3.1 by
changing the densities of six elements involved in the formation of the
observed emission lines: nitrogen, oxygen, neon, sulphur, argon, and iron.
These six ions do not independently act but are tightly correlated to each
other. For example, the increased neon abundance does not always lead to
stronger neon lines, or the increased iron abundance strengthens the [Ne II]
line as well as iron lines but not the [Ne III] line. Since hydrogen is
depleted in IRS1, heating process mostly depends on photoelectric heating by
dust and heavy elements rather than photoionization by hydrogen; thus, changes
of the metal and dust abundances affect the heating and cooling processes,
complicating the model calculations. With a large number of free parameters
and limited observational data, it is unlikely possible to find the only model
that perfectly fits IRS1. Instead, we intend to present a reference model that
reasonably well explains the observations and to discuss some parameters that
affect the modeling results.

In Figure 8, we present the reference model of IRS1 in red color that fits
fairly well the Spitzer IRS spectrum in both dust features and line
intensities. The model spectrum was smoothed to the spectral resolution of the
Spitzer IRS modules. The model reproduces the observed dust features at 18,
23, 27, and 34 µm but with larger flux and narrower width for the 23 µm peak.
The continuum slope from 15 to 20 µm in the model is also steeper than
observed. These differences can be explained by the sizes and shapes of dust
grains. We have assumed a spherical dust grain of 0.1 µm (or 0.25 µm for FeO)
because of the limited availability of the optical constants or opacities of
dust in MIR, but dust properties in fact highly depend on both dust size and
shape (e.g., Koike et al., 1989, 2010; Min, 2015). For example, the 10 µm
silicate feature becomes broader and is shifted to longer wavelength as the
grain shape deviates from a perfect sphere. The prominent features of
forsterite (Mg-rich crystalline silicate) are also significantly suppressed
for the dust grains with larger size (Figure 5 of Min, 2015) or elliptical
shape (Figure 9 of Koike et al., 2010). This implies that the model of the IRS
spectrum can be improved by using elliptical and/or larger dust grains. We
note that the dust models in Koo et al. (2011) that used a continuous
distribution of ellipsoids (CDE) for FeO give a better fit for the steep
continuum shape at 15–20 µm, but their models produced by independent dust
components result in the temperature of 90–150 K for FeO and Mg0.6Fe0.4O (or
MgO), which is much higher than the temperature of the other dust components
around 55 K.

Figure 8: Spitzer IRS spectrum (black) of IRS1 and Cloudy models with different thickness. The red line is the reference model that fits fairly well the Spitzer IRS spectrum in both dust features and line intensities. The purple and blue lines are the models calculated with the same parameters as the reference model but with a half and twice of the thickness. The model spectra were smoothed to the spectral resolution of the Spitzer IRS modules.  Table 4: Dust Parameters of IRS1 from the Reference Model  Component | Mass | Temperature
---|---|---
| ($10^{-3}~{}M_{\sun}$) | (K)
Crystalline olivine (Mg1.9Fe0.1SiO4) | 0.06 | 79
FeO | 0.04 | 67
Mg0.6Fe0.4O | 0.20 | 79
Amorphous silicate | 2.09 | 72
Total | 2.40 |  $\cdots$
Table 5: Line Fluxes Predicted from the CLOUDY Models and Relative Intensities
to the Observations

ModelaaA scale factor applied to the cloud thickness. | [Ar III] | [S IV] | [Ne II] | [Ne III] | [S III] | [Fe III]bbNot seen in the Spitzer IRS spectrum, but likely blended with dust features. | [Fe II] | [S III]bbNot seen in the Spitzer IRS spectrum, but likely blended with dust features.
---|---|---|---|---|---|---|---|---
($\times$ thick) | 8.99 µm | 10.51 µm | 12.81 µm | 15.56 µm | 18.71 µm | 22.93 µm | 25.99 µm | 33.48 µm
0.5 | 1.07E-13 (0.80) | 1.22E-13 (0.87) | 7.66E-12 (0.93) | 9.03E-13 (0.93) | 3.48E-13 (0.68) | 2.68E-12 | 7.35E-15 (0.01) | 3.62E-12
1.0 | 1.31E-13 (0.99) | 1.27E-13 (0.90) | 9.16E-12 (1.12) | 9.08E-13 (0.94) | 5.45E-13 (1.07) | 2.31E-11 | 1.84E-13 (0.20) | 6.60E-12
2.0 | 1.42E-13 (1.07) | 1.29E-13 (0.91) | 1.01E-11 (1.23) | 9.10E-13 (0.94) | 6.97E-13 (1.37) | 5.35E-11 | 7.60E-13 (0.82) | 9.50E-12

Note. — Line flux is in erg s-1 cm-2. The numbers in parentheses are the
relative intensities with respect to the observed fluxes.

Figure 9:  Ionization fraction of neon (left) and iron (right) from the Cloudy
models of IRS1, which is defined as the number density of each ion of an
element among the total number density of the element. The solid lines are the
reference model; the dashed and dotted lines are the models with twice and a
half of the thickness, respectively. Ions are presented by different colors:
navy (Ne II, Fe II), blue (Ne III, Fe III), and cyan (Fe IV).

Table 4 presents dust mass and temperature obtained from the reference model.
Most dust mass is contributed by amorphous silicate. The contribution from
crystalline olivine is small but indispensable to fit the observed dust
features. The total dust mass is much smaller ($\sim$27%) than the dust mass
of $9\times 10^{-3}$ $M_{\sun}$ derived from the modified-blackbody fit (Koo
et al., 2011), likely due to the constrained geometry. Total dust mass highly
depends on the geometry such as $r_{0}$ or the thickness of IRS1 both of which
were fixed in our calculations. If we increase $r_{0}$ to 0.52 pc, dust mass
becomes comparable to that of Koo et al. (2011). We note that dust mass from
our calculation is consistent with dust mass from the model with a single dust
of carbonaceous-silicate (Koo et al., 2011). Comparing to Koo et al. (2011),
the relative fraction and temperature of the individual dust species are also
different. This may come from the differences in dust opacity. While the
optical constants of each dust were adopted from the same literature, the
calculations of dust absorption coefficients are different likely due to the
assumed shapes and sizes of dust grain, resulting in slightly different dust
properties. For example, the strength of two peaks at 18 µm and 23 µm of the
absorption coefficient of crystalline olivine are comparable in our
calculations (Figure 6), while the 23 µm peak is only 60% relative to the 18
µm peak in Koo et al. (2011), requiring larger mass and lower temperature to
fit the observed 23 µm feature. For FeO, we assume a spherical grain of $0.25$
µm, whereas Koo et al. (2011) assumed a CDE with a size of 0.1 µm which shows
the weaker, broader, and asymmetric peak of the absorption coefficient.

The Cloudy model with the abundance listed in Table 3 also reproduces several
ionic lines observed in the IRS spectrum. Since we have applied a scale factor
of 7.5 in log scale to metals and grains in the abundance of the nova V1500
Cyg and adjusted the amounts of selected ions, we also present the relative
abundance of metals to the carbon abundance. The reference Cloudy model
indicates that neon, argon, and iron are enhanced in IRS1. As we pointed out
earlier, the reference model is not the only model that can explain IRS1, but
the overall trend in abundance may not be very different in order to produce
the observed lines. The predicted lines and their fluxes are presented in
Table 5 with the relative strength to the observed line fluxes. The predicted
line fluxes mostly agree with the observations. The [S III] line at 33.48 µm
and [Fe III] line at 22.93 µm are also predicted. While these two lines are
predicted to be strong, they are blended with the dust features of crystalline
olivine at 34 µm and 23 µm, respectively, so hardly detectable in the IRS
spectrum. In the model spectrum, after we smoothed it to the Spitzer IRS
resolution, the [S III]33.48µm and the [Fe III]22.93µm lines have become
hidden under the dust features as well (Figure 8).

On the other hand, the [Fe II]25.99µm line predicted from the model is much
weaker than observed. While we noted that the line detected at $\sim$25.9 µm
could be either [O IV]25.89µm or [Fe II]25.99µm, the Cloudy model only
predicts [Fe II]25.99µm. We also believe that [Fe II] is more plausible
because the UV radiation of an O4.5III star is not hard enough to ionize O III
to O IV of which ionization potential is 54.9 eV. The reason for the weaker
[Fe II] line in the model is because most iron is in Fe III or Fe IV as shown
in Figure 9 (see also Bautista & Pradhan, 1998). The [Fe II] line becomes
stronger when the thickness of IRS1 increases. In Figure 8 and Table 5, we
present two models produced by the same parameters as the reference model
except for the cloud thickness: the models with the thickness twice and a half
of the thickness of the reference model. In Table 5, the [Fe II]25.99µm and
[Fe III]22.93µm lines are the only lines that significantly vary by thickness.
This behavior is also seen in Figure 9 that presents the ionization fraction
of neon and iron by depth. In the figure, depth is normalized and is zero at
the closest side to the heating source, where the temperature is the highest.
The ionization fraction is defined as the number density of each ion of an
element among the total number density of the element. Figure 9 shows that the
ionization fraction of iron is sensitive to the thickness compared to neon.
For example, all of the neon is in Ne II from the very inside ($\gtrsim$ 0.15
of the depth) regardless of the cloud thickness. In contrast, the fraction of
iron ions changes through the whole depth depending on the cloud thickness.
Therefore, the model of IRS1 could be improved by leaving the cloud thickness
as a free parameter or by finding a constraint of the thickness that
reproduces the observed [Fe II]25.99µm line.

In Cloudy model calculations, we have assumed a lack of hydrogen based on the
SED models (Section 3.4). We now examine a possibility of a higher hydrogen
density. In Figure 10, we present the reference model without hydrogen, i.e.,
$n({\rm H})=10^{-4.9}$ cm-3, and another two models including some amount of
hydrogen, i.e., $n({\rm H})=10$ and 100 cm-3, with the observed SED of IRS1.
The MIR spectra of the three models are identical as long as the total amount
of metals and grains are retained the same, but the models with hydrogen show
a group of emission lines from hydrogen in optical and NIR which have not been
observed. We also searched for the H$\alpha$ emission around IRS1 from the VST
Photometric H$\alpha$ Survey of the Southern Galactic Plane and Bulge (VPHAS+;
Drew et al., 2014)888http://www.vphasplus.org. The predicted H$\alpha$ line
flux from the three models of IRS1 are $4.2\times 10^{-21}$, $4.2\times
10^{-15}$, and $1\times 10^{-13}$ erg s-1 cm-2 for the models with $n({\rm
H})=10^{-4.9}$, 10, and 100 cm-3, respectively, after applying the extinction
of $A_{\rm V}\sim 5$ mag (from $N_{\rm H}=9.2\times 10^{21}$ cm-2; Koo et al.,
2011). For comparison, the 5$\sigma$ limiting magnitude of VPHAS+ is about 20
mag (Drew et al., 2014), or $1.84\times 10^{-15}$ erg s-1 cm-2 in H$\alpha$.
If IRS1 contains hydrogen, even a small amount of $\lesssim 10$ cm-3, the
H$\alpha$ emission is expected to be detected in the VPHAS+ images, but no
emission has been found around IRS1. This supports our assumption that
hydrogen is depleted in IRS1. We note that there is still weak continuum
emission at short wavelength in the reference model without hydrogen, a little
stronger than the upper limits from the observation in optical and NIR. This
continuum is the reflected emission of the incident radiation field from the
heating source, so it may not be observed if it is diffuse.

Figure 10: SEDs of IRS1 from the observations and Cloudy models with different
hydrogen density. The red (hydrogen-deficient; $n({\rm H})=10^{-4.9}$ cm-3)
line is the reference model the same as the red line in Figure 8. The blue
($n({\rm H})=10$ cm-3) and purple ($n({\rm H})=10^{2}$ cm-3) lines are the
models with hydrogen.

##  4 Crystalline Silicate in MSH 15$-$52

MSH 15$-$52 is the first SNR in which crystalline silicate is observed. Our
analysis in this paper indicates that the elemental abundance of the IR
compact source IRS1 where crystalline silicate has been detected is close to
that of SN ejecta with depleted hydrogen and high abundance of metals,
particularly neon, argon, and iron. This implies that IRS1 (and probably IRAS
15099-5856 as well) originates from the SN ejecta rather than the mass loss of
the SN progenitor as has been proposed by Koo et al. (2011). If this is true,
MSH 15$-$52, besides the existence of crystalline silicate, is a very unique
object where we can directly observe dust newly formed in the ejecta of SNe
Ib/c, which would have not been possible without the nearby heating source
Muzzio 10. While FIR observations (e.g., AKARI, Herschel, and ISO) have
revealed cold dust inside SNRs, the spatial resolutions are not high enough to
disentangle the SN dust from the surrounding ISM dust and examine dust
properties in detail (e.g., Sibthorpe et al., 2010; Koo et al., 2016; Chawner
et al., 2020; Millard et al., 2021; Rho et al., 2023, and references therein).
The cold and warm dust has been detected in MSH 15$-$52 as well. Millard et
al. (2021) estimated 0.03–0.06 $M_{\sun}$ of warm (46–52 K) and 4–15
$M_{\sun}$ of cold (17–20 K) dust, assuming the distance of 5.2 kpc to MSH
15$-$52, from the two-component blackbody model fitting of the MIR to FIR
spectrum obtained by the Long Wavelength Spectrometer (LWS) on board the ISO.
Since their spectrum is not background-subtracted, they suggest that the warm
and cold dust originate from the SN ejecta and background ISM, respectively.
The dust mass of 0.03–0.06 $M_{\sun}$ at 5.2 kpc is scaled to 0.018–0.036
$M_{\sun}$ at 4 kpc. This is ten times larger than our results, but a direct
comparison of the two dust masses is inappropriate because the LWS spectrum of
MSH 15$-$52 with a large beam size of $\sim$80$\arcsec$ (Gry et al., 2003)
includes not only IRS1 but the surrounding, diffuse emission.

Dust formation in the SN ejecta of Type II SNe (SNe II) has been widely
studied by theoretical calculations (e.g., Sarangi & Cherchneff, 2013, 2015;
Brooker et al., 2022) as well as the observations of young SNRs, e.g., Cas A
(De Looze et al., 2017), SN 1987A (Matsuura et al., 2015), G54.1+0.3 (Rho et
al., 2018), and the Crab Nebula (Gomez et al., 2012). In contrast, few study
thus far has been carried out for dust condensation in the ejecta of SNe Ib/c.
Observational signatures of dust formation in the SN Ib/c ejecta have been
discovered only for a few SNe in the nebula phase: SN Ib 1990I (Elmhamdi et
al., 2004), SN Ib(n) 2006jc (Di Carlo et al., 2008; Smith et al., 2008), SN Ic
2020oi (Rho et al., 2021), and SN Ic 2021krf (Ravi et al., 2023). The
molecules which become dust seeds in the SN ejecta and their chemical
reactions would not be very different between SNe II and Ib/c, but the
environments in which dust condensation occurs would differ depending on SN
types and progenitor stars. For example, even in the same SN type of IIP, the
amount of dust formed in the ejecta and the degree of grain growth predicted
by dust condensation models highly depend on the conditions of SN explosion
such as progenitor mass, explosion energy, mass of 56Ni, or clumpy structure
of ejecta (Sarangi & Cherchneff, 2013, 2015; Brooker et al., 2022). The
observations of dust signatures in the SN Ib/c ejecta listed above indicate
that the onset of dust formation is much earlier for SNe Ib/c that is found to
be 50–70 days after SN explosion (Di Carlo et al., 2008; Smith et al., 2008;
Rho et al., 2021; Ravi et al., 2023) except SN 1990I with 230 days (Elmhamdi
et al., 2004) than for SNe II that is estimated to be later than a few
hundreds days after explosion (Sarangi et al., 2018, and references therein).
The reason for the early dust condensation in the ejecta of SNe Ib/c is
thought to be due to the rapid decrease in the gas temperature (Nozawa et al.,
2008). For SNe Ib/c, the ejected masses are smaller and the expansion
velocities are higher than SNe II because the SN Ib/c progenitors have lost
most of the hydrogen/helium envelopes before explosion. This leads to a lower
density of gas in the ejecta, and gas temperature drops down more quickly than
those in typical SNe II. Nozawa et al. (2008) calculated dust formation in the
SN 2006jc applying the SN Ib model of a relatively low-mass (6.9 $M_{\sun}$)
helium star progenitor with an ejecta mass of 4.9 $M_{\sun}$. Their
calculation predicts the gas temperature reaching a typical dust condensation
temperature ranges of 1,000–2,000 K between 50 and 200 days after the SN
explosion, which means that dust forms much closer to the explosion center in
the ejecta of SNe Ib/c than in the SNe II ejecta (Figure 1 of Nozawa et al.,
2008). A similar process of dust formation might have occurred in MSH 15$-$52
of which progenitor is also speculated as a low-mass helium star.

The different environmental conditions may bring distinctive characteristics
of dust formed in the SN Ib/c ejecta, for example, the formation of
crystalline silicate or crystallization of amorphous silicate. If silicates
are formed at high ($>$1,000 K) temperature, the crystalline lattice structure
is the most favorable state (Molster & Kemper, 2005). In a clumpy ejecta,
stoichiometric silicate (Mg2SiO4) is predicted to be formed at high densities
(Sarangi & Cherchneff, 2015). Therefore, the condensation of crystalline
silicates could take place in the SN ejecta, or silicates first formed in
amorphous structure could be crystallized if there is a high-energy process
such as heating by pulsar wind nebula, although the condensation temperature
and condensing phase of dust are not simply determined but associated with
several factors such as gas pressure or gas kinematics (Nagahara et al., 2009;
Gail et al., 2013). Then, assuming that crystalline silicates can form in SN
dust, why have crystalline silicates not been found in other SNRs except MSH
15$-$52? This is probably because it is difficult to observe unshocked SN dust
without a heating source such as Muzzio 10. Crystalline silicates, if present,
can be undetected in FIR because they exhibit very weak or no spectral
signatures at wavelengths longer than 40 µm except for the 69 µm feature
(Koike et al., 2003; Sturm et al., 2013).

While our results suggest a possibility that the crystalline silicate of IRS1
originates from the SN ejecta, the current observational data with limited
spatial and spectral resolution can neither confirm the ejecta origin nor rule
out the progenitor origin. The Gemini/T-ReCS images show a slightly different
spatial distribution between [Ne II] 12.81 µm and Qa 18.30 µm in a spatial
scale less than one arcsecond (Figure 2). This is likely because the [Ne II]
line does not trace the Mg silicate but dust with a very smooth spectrum such
as Al2O3 (Arendt et al., 2014). The spatial distributions of various ionic
lines will be required to examine a correlation between the gas and dust
features associated with crystalline silicate. The Spitzer IRS spectrum of
IRS1 did not resolve the velocity of the gas component. Previously, Koo et al.
(2011) favored the progenitor origin of IRS1 based on the low ($-160\pm 560$
km s-1) central velocity of the [Ne II] line. The velocities of the other
lines except [Fe II] 25.99 µm are similarly a few hundreds km s-1 with an
average of $-444$ km s-1, but the uncertainties are huge as [Ne II]. The [Fe
II] line exceptionally shows a large velocity of $-1710\pm 280$ km s-1. This
may imply a different origin of [Fe II] from the other lines, but it can be
due to an inaccurate measurement of the velocity since the [Fe II] line is
weak and embedded between strong dust features (Figure 3).

To confirm the origin of the crystalline silicate in MSH 15$-$52 and explore
the possibility of the formation of crystalline silicates in SN ejecta,
further observations as well as theoretical investigations are required.
Particularly, it is crucial to examine the spatial distributions of gas and
dust in IRAS 15099-5856 through the MIR observations with high resolution and
sensitivity (e.g., JWST/MIRI). If the origin of the crystalline silicate is
found to be the SN ejecta, MSH 15$-$52 will give an unprecedented opportunity
to investigate dust formation in SN ejecta, which is not yet clearly known.

##  5 Summary and Conclusion

We have presented the MIR imaging observations and analysis of the compact IR
source IRS1 of IRAS 15099-5856 in the SNR MSH 15$-$52, which is the first and
only object with crystalline silicate dust associated with SNRs so far. The
MIR images obtained by using Gemini/T ReCS revealed the morphology of IRS1 and
spatial distributions of gas and dust at a spatial resolution of $\lesssim$
1$\arcsec$. We have also presented the analysis of the Spitzer IRS spectrum of
IRS1 that was previously investigated with the models of thermal emission from
multiple independent dust components (Koo et al., 2011). In this paper, we
have analyzed the ionic lines and modeled the spectrum considering the
geometry and energy balance to derive the chemical abundance of gas as well as
dust parameters. The derived abundance is close to that of SN ejecta with poor
hydrogen and enhanced metals. This suggests the ejecta origin for the
crystalline silicate and may imply the possibility of the formation of
crystalline silicate in SN ejecta, but the current observational data are
still limited in spatial and spectral resolution. If the origin of the
crystalline silicate in IRAS 15099-5856 is confirmed as the SN ejecta by
future observations, MSH 15$-$52 will be a very unique, invaluable object that
proves the formation of crystalline silicate in SN ejecta and where we can
directly observe newly-formed dust in the ejecta of SNe Ib/c. In the
following, we summarize our main results.

  1. 1.

The Gemini/T-ReCS images show a complicated, extended morphology of IRS1 with
bright clumps and diffuse emission in [Ne II] 12.81 µm and Qa 18.30 µm. The
[Ne II]cont image with no emission and the Si-6 image with the almost same
emission as [Ne II] indicate that there is no other line or strong continuum
emission. The T-ReCS images confirm the previous prediction (Koo et al., 2011)
that IRS1 is extended and externally heated by the nearby O star Muzzio 10. We
estimated the [Ne II] and Qa flux of IRS1 from the T-ReCS images and compared
them with the flux derived from the Spitzer IRS spectrum.

  2. 2.

The Spitzer spectrum of IRS1 shows prominent dust features at 23, 27, and 34
µm that can be explained by crystalline silicate dust. We also detected
several ionic lines of [Ar III] 8.99 µm, [S IV] 10.51 µm, [Ne II] 12.81 µm,
[Ne III] 15.56 µm, [S III] 18.71 µm, and [O IV] 25.89/[Fe II] 25.99 µm. The [O
IV] and [Fe II] are not resolved at the spectral resolving power of the
Spitzer IRS LL module, but it is likely [Fe II]. The estimated line flux is
from 0.13 to 8.21 $\times 10^{-12}$ erg s-1 cm-2. The line widths are
comparable to the spectral resolving power, i.e., the velocity is not
resolved.

  3. 3.

We compared the line ratios of [Ne III]15.56µm/[Ne II]12.81µm versus [S
IV]10.51µm/[S III]18.71µm of various astronomical objects with the observed
line ratios of IRS1 on a model grid generated by Cloudy (Ferland et al.,
2013). The line ratio diagram shows that the abundance of IRS1 is rather close
to the nova abundance with enhanced neon. The absence of hydrogen line in the
Spitzer spectrum further suggests that hydrogen is depleted in IRS1.

  4. 4.

We modeled the Spitzer spectrum of IRS1 using the photoionization code Cloudy
(Ferland et al., 2013). We assumed the cloud IRS1 externally heated by an
O4.5III star (Muzzio 10) separated by 0.45 pc. For the gas, the nova abundance
was initially adopted but with hydrogen depleted. For the dust species,
crystalline olivine (Mg1.9Fe0.1SiO4), FeO, Mg0.6Fe0.4O, and amorphous silicate
were included. Spherical dust grains of 0.25 µm (FeO) and 0.1 µm (the others)
were assumed. We first fitted the dust features and adjusted the abundance of
nitrogen, oxygen, neon, sulphur, argon, and iron to find a model that
reproduces the observed lines. We have derived a reference model that fairly
well fits the Spitzer spectrum and discussed the factors that affect the
models.

  5. 5.

The reference model fits the dust features at 27 and 34 µm, while it does not
well fit the 23 µm feature and steeply-increasing continuum between 15 and 20
µm. The derived dust mass is $2.4\times 10^{-3}$ $M_{\sun}$ in total mostly
contributed by amorphous silicate. The dust temperature is between 67 and 79
K. The total dust mass is about 30% of the mass previously derived by multi-
component blackbody fitting (Koo et al., 2011), but the dust mass depends on
the separation to the heating source which was fixed in our modeling. The
reference model predicts the emission lines with comparable intensities to the
observation except the [Fe II] 25.99 µm line which is predicted much weaker
(20%) than the observed intensity. We have found that the ionization fraction
of iron is sensitive to the cloud thickness which was also fixed in our
modeling. The model can be improved by changing dust properties (shape or
size) and geometrical parameters (e.g, cloud thickness or the separation
between IRS1 and Muzzio 10), although it is beyond the scope of this paper.

  6. 6.

The Cloudy modeling of the Spitzer spectrum indicates that the elemental
abundance of IRS1 is close to that of SN ejecta with depleted hydrogen and
high abundance of metals, particularly neon, argon, and iron. This implies
that IRS1 (and IRAS 15099-5856) is the SN ejecta and that crystalline silicate
can form in SN ejecta. To confirm this, further observations to resolve the
spatial distributions of various ionic lines, dust continuum, and crystalline
silicate will be required.

This work is based in part on observations obtained at the international
Gemini Observatory, a program of NSF’s NOIRLab, which is managed by the
Association of Universities for Research in Astronomy (AURA) under a
cooperative agreement with the National Science Foundation on behalf of the
Gemini Observatory partnership: the National Science Foundation (United
States), National Research Council (Canada), Agencia Nacional de Investigación
y Desarrollo (Chile), Ministerio de Ciencia, Tecnología e Innovación
(Argentina), Ministério da Ciência, Tecnologia, Inovações e Comunicações
(Brazil), and Korea Astronomy and Space Science Institute (Republic of Korea).
The Gemini program identification number associated with the data is
GS-2012A-C-4. This work is based in part on observations made with the Spitzer
Space Telescope, which was operated by the Jet Propulsion Laboratory,
California Institute of Technology under a contract with NASA. This work has
made use of data from the European Space Agency (ESA) mission Gaia
(https://www.cosmos.esa.int/gaia), processed by the Gaia Data Processing and
Analysis Consortium (DPAC,
https://www.cosmos.esa.int/web/gaia/dpac/consortium). Funding for the DPAC has
been provided by national institutions, in particular the institutions
participating in the Gaia Multilateral Agreement. This research of H.-J. K.
was supported by Basic Science Research Program through the National Research
Foundation of Korea (NRF) funded by the Ministry of Education
(RS-2023-00246733). The work of B.-C. K. was supported by the Basic Science
Research Program through the National Research Foundation of Korea (NRF)
funded by the Ministry of Science, ICT and Future Planning
(2020R1A2B5B01001994, RS-2023-00277370).

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