The Astrophysical Journal, 706:441–453, 2009 November 20

doi:10.1088/0004-637X/706/1/441


C 2009. The American Astronomical Society. All rights reserved. Printed in the U.S.A.

AKARI INFRARED OBSERVATIONS OF THE SUPERNOVA REMNANT G292.0+1.8: UNVEILINGCIRCUMSTELLAR MEDIUM AND SUPERNOVA EJECTA
Ho-Gyu Lee1,2 , Bon-Chul Koo2 , Dae-Sik Moon3 , Itsuki Sakon1 , Takashi Onaka1 , Woong-Seob Jeong4 ,Hidehiro Kaneda5 , Takaya Nozawa6 , and Takashi Kozasa7
1 Department of Astronomy, Graduate School of Science, The University of Tokyo, Bunkyo-ku, Tokyo 113-0033, Japan; hglee@astron.s.u-tokyo.ac.jp,

isakon@astron.s.u-tokyo.ac.jp, onaka@astron.s.u-tokyo.ac.jp
2 Department of Physics and Astronomy, Seoul National University, Seoul 151-742, Republic of Korea; koo@astrohi.snu.ac.kr
3 Department of Astronomy and Astrophysics, University of Toronto, Toronto, ON M5S 3H4, Canada; moon@astro.utoronto.ca
4 Korea Astronomy and Space Science Institute, 61-1, Whaam-dong, Yuseong-gu, Deajeon 305-348, Republic of Korea; jeongws@kasi.re.kr
5 Graduate School of Science, Nagoya University, Chikusa-ku, Nagoya 464-8602, Japan; kaneda@u.phys.nagoya-u.ac.jp
6 Institute for the Physics and Mathematics of the Universe, University of Tokyo, Kashiwa, Chiba 277-8568, Japan; tnozawa@mail.sci.hokudai.ac.jp
7 Department of Cosmosciences, Graduate School of Science, Hokkaido University, Sapporo 060-0810, Japan; kozasa@mail.sci.hokudai.ac.jp

Received 2009 May 7; accepted 2009 October 7; published 2009 October 30

ABSTRACT
We present the results of AKARI observations of the O-rich supernova remnant (SNR) G292.0+1.8 using six
Infrared Camera (IRC) and four Far-Infrared Surveyor bands covering 2.7–26.5 μm and 50–180 μm, respectively.
The AKARI images show two prominent structures; a bright equatorial ring (ER) structure along the east–west
direction and an outer elliptical shell structure surrounding the remnant. The ER structure is clumpy and incompletewith its western end opened. The outer shell is almost complete and slightly squeezed along the north–south
direction. The central position of the outer shell is ∼1 northwest from the embedded pulsar and coincides withthe center of the ER structure. In the northern and southwestern regions, there is also faint emission with a sharpboundary beyond the bright shell structure. The ER and the elliptical shell structures were partly visible in optical
and/or X-rays, but they are much more clearly revealed in our AKARI images. There is no evident difference
in infrared colors of the two prominent structures, which is consistent with the previous proposition that both
structures are of circumstellar origin. However, we have detected faint infrared emission of a considerably high
15/24 μm ratio associated with the supernova (SN) ejecta in the southeastern and northwestern areas. Our IRC
spectra show that the high ratio is at least partly due to the emission lines from Ne ions in the SN ejecta material.
In addition, we detect a narrow, elongated feature outside the SNR shell. We derive the physical parametersof the infrared-emitting dust grains in the shocked circumstellar medium (CSM) and compare the result withmodel calculations of dust destruction by an SN shock. The AKARI results suggest that the progenitor was at
the center of the infrared circumstellar shell in the red supergiant stage and that the observed asymmetry in
the SN ejecta could be a result of either a dense CSM in the equatorial plane and/or an asymmetric explosion.
Key words: infrared: ISM – ISM: individual (G292.0+1.8) – shock waves – supernova remnants

kinematics of the [O iii] line emission from the ejecta, is roughly
coincident with the geometrical center of the radio emission,
but shows an apparent offset from its pulsar position discoveredat the southeast (Winkler et al. 2009, see below). Ghavamian
et al. (2009) reported that the mid-infrared (MIR) spectrumof the ejecta is dominated by ionic lines and a broad bumparound 17 μm using the Spitzer Space Telescope (Spitzer).
They proposed that the latter is produced either by polycyclicaromatic hydrocarbons (PAHs) along the line of sight or newly
formed dust within the ejecta. Besides, there is also an equatorial
MIR continuum component from the dust associated with theshocked CSM partly overlapping with the ejecta emission.
A previous infrared (IR) study based on the data obtainedwith Infrared Astronomical Satellite (IRAS), on the other hand,
identified enhanced far-IR (FIR) emission in the southwesternpart of the SNR, suggesting that the source has encounteredadjacent clouds (Braun et al. 1986; Park et al. 2007).
In radio, G292.0+1.8 mainly consists of bright emission froma central pulsar wind nebula (PWN) of ∼4 (in diameter) andrelatively fainter outer plateau emission of ∼9 (Lockhart et al.
1977; Braun et al. 1986; Gaensler & Wallace 2003). The plateau
emission declines sharply outward, and there is no apparent
surrounding radio shell structure. The pulsar PSR J1124–5916,
which is located ∼46 from the dynamical center of the SNR in

1. INTRODUCTION
Young core-collapse supernova remnants (SNRs) in the
Galaxy provide an unique opportunity to study fine detailsof ejecta from supernova (SN) explosions as well as thoseof the circumstellar medium (CSM) produced over the final
evolutionary stages of massive stars. An obvious group of suchyoung core-collapse SNRs is O-rich SNRs that show optical
spectra featuring strong O and Ne lines, with lines from lighter
elements (e.g., He) either weak or absent. Since the ejecta froma progenitor and the swept-up CSM in the O-rich SNRs are
not completely mixed together yet, observation of the O-rich
SNRs is promising to study the ejecta and CSM related to the
explosion of a progenitor and its late-stage evolution.
G292.0+1.8 (MSH11–54), together with Cassiopeia A and
Puppis A, forms a rare group of O-rich SNRs in the Galaxy.
The O-rich nature of G292.0+1.8 was discovered by the opticaldetection of fast-moving O-rich and Ne-rich ejecta (Goss et al.
1979; Murdin & Clark 1979). Recent wide-field observationsof the optical [O iii] line in G292.0+1.8 have revealed that theejecta velocities range from −1400 km s−1 to +1700 km s−1
and also that distant ejecta knots are located primarily along
the north–south direction (Ghavamian et al. 2005; Winkler &Long 2006). The dynamical center of G292.0+1.8, based on the
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Table 1
Summary of AKARI Observations
Date

IDa

Mode

Band

2007 Jan 172007 Jan 172007 Jan 192007 Jul 21
2007 Jul 20
2007 Jul 21–22

1400748
1400749
1400751
1402801
1402839, 41, 43
1402840, 42, 44

IRC imaging
IRC imaging
FIS imaging
IRC imaging
IRC spectroscopy
IRC spectroscopy

N3, N4, S7, S11
L15, L24
N60, Wide-S, Wide-L, N160
L18W
NIR, MIR-S
MIR-L

Note.
a AKARI observational identification number. These observations are part of the AKARI mission program, ISM in
our Galaxy and nearby galaxies (PI: H. Kaneda).
Table 2
Characteristics of AKARI Imaging Observations
Instrument
ChannelBand
Reference wavelength (μm)
Effective bandwidth (μm)
Pixel size ( )FWHM ( )IRC
NIR
512 × 412 InSb

MIR-S
256 × 256 Si:As

N3
3.2
0.87
1.46
4.0

S7
7.0
1.75
2.34
5.1

N4
4.1
1.53
1.46
4.2

S11
11.0
4.12
2.34
4.8

FIS
MIR-L
256 × 256 Si:As
L15
15.0
5.98
2.51a
5.7

L24
24.0
5.34
2.51a
6.8

SW
Ge:Ga
N60
65
21.7
26.8
37

Wide-S
90
37.9
26.8
39

LW
Stressed Ge:Ga
Wide-L
140
52.4
44.2
58

N160
160
34.1
44.2
61

Note. a Represents the width of a pixel. The exact size is 2. 51 × 2. 34.

the southeast direction, was discovered by radio timing (Camilo
et al. 2002), confirming the PWN nature of the central emissionof the SNR. The pulsar and PWN were also identified in X-ray
emission (Hughes et al. 2003), where the pulsar has a nearby jet
and a torus of ∼5 scale (Park et al. 2007). The jet axis is slightly
tilted to the northeast direction (Park et al. 2007). There are twoclear differences between the soft and hard X-ray emission of
G292.0+1.8. First, while the former is produced by shockedgas of normal composition distributed as equatorial filaments,
indicating it is likely from CSM, the latter is dominated bybright metallic lines from the ejecta (Park et al. 2002, 2004).
Next, the X-ray emission is harder in the northwestern part than
the southeastern part, implying density variation or asymmetric
SN explosion (Park et al. 2007). The distance to G292.0+1.8, on
the other hand, was estimated to be 6.2 kpc from H i absorption
observations (Gaensler & Wallace 2003).
In this paper, we present an extensive IR study of G292.0+1.8
using multi-band imaging and spectroscopic data obtained with
the AKARI space telescope (Murakami et al 2007). We describe
the details of our AKARI observations and present the resultsof observations in Sections 2 and 3, respectively. In Section 4,
we discuss IR emission in G292.0+1.8 from circumstellar dustwith a comparison to model calculations of dust destruction by a
SN shock. We also discuss the IR emission from the ejecta, andwe suggest that there is a negligible amount of dust associatedwith the ejecta. Then, we make a comparison to the results ofSpitzer observations. We finally discuss the SN explosion in
G292.0+1.8 based on our AKARI observations followed by ourconclusions in Section 5.
2. OBSERVATIONSThe AKARI multi-band imaging observations of G292.0+1.8
were carried out using its six near-IR (NIR) and MIR bands
in the 2.7–26.5 μm range as well as four FIR bands in the
50–180 μm range on 2007 January 17 and 19. The NIR andMIR images were obtained with a Infrared Camera (IRC) thathas three simultaneously operating NIR, MIR-S, and MIR-L
channels. The three channels have comparable fields of view
(FoVs) of ∼10 × 10 . The NIR and MIR-S channels share thesame pointing direction, while that of the MIR-L channel is25 away. The IRC observations were conducted in the twofilter mode that produced images of two bands for each channel(Onaka et al. 2007). The total on-source integration times were
178 s for the NIR (N3, N4) observations, and 196 s for both
the MIR-S (S7, S11) and MIR-L (L15, L24) observations. The
basic calibration and data reduction, including dark subtraction,
linearity correction, distortion correction, flat fielding, image
stacking, and absolute position determination were performedby the standard IRC Imaging Data Reduction Pipeline (version
20071017).8 Tables 1 and 2 present a journal of our AKARI
observations, including the spectroscopic observations, and thebasic parameters of the AKARI imaging bands, respectively.
The FIR images were obtained with Far-Infrared Surveyor(FIS) in two round-trip scans using the cross-scan shift mode
(Kawada et al. 2007). The scan speed and length were 15 s−1
and 240 , respectively, creating images of 40 × 12 size
elongated along the scanning direction. All the four FIS band
(N60, Wide-S, Wide-L, and N160) images were obtained
simultaneously by a single observing run (Table 1). The initial
data calibration and reduction such as glitch detection, dark
subtraction, flat fielding, and flux calibration were processedwith FIS Slow-Scan Toolkit (version 20070914),9 and the final
image construction was performed with a refined samplingmechanism.
These IRC and FIS multi-band imaging observations were
followed by NIR and MIR spectroscopic observations, andL18W-band (13.9–25.6 μm, combining both L15 and L24
bands) imaging observations carried out in 2007 July 20–22.
The spectroscopic observations were conducted using four (NG,89

http://www.ir.isas.jaxa.jp/ASTRO-F/Observation/DataReduction/IRC/
http://www.ir.isas.jaxa.jp/ASTRO-F/Observation/DataReduction/FIS/

No. 1, 2009AKARI INFRARED OBSERVATIONS OF THE SNR G292.0+1.8

Table 3Characteristics of AKARI Spectroscopic Observations
Channel

NIR

Grism
Slit size ( × )Wavelength coverage (μm)
Resolving power
Exposure time (s)NG5 × 0. 8
2.5–5.0
120
400

MIR-S
SG1
5 × 0. 8
4.6–9.2
53
196

MIR-L

SG2
5 × 0. 8
7.2–13.4
50
245

LG2
7 × 0. 4
17.5–26.5
48
442

SG1, SG2, and LG2) grisms. Similar to the aforementioned
imaging observations, the NG, SG1, and SG2 mode observations
were conducted simultaneously, while LG2 mode observations
were done separately (Ohyama et al. 2007). Table 3 lists
the characteristics of the AKARI spectroscopic observations.
Spectra from the peaks of the equatorial emission and ejecta
identified in the ratio between the L15- and L24-band images(L15/L24 hereafter) were obtained together with that froma reference background position at the southeastern part ofthe source (Table 4). The data calibration and reduction were
processed with IRC Spectroscopy Toolkit (version 20070913).10
The obtained flux was converted to the surface brightness basedon the measured slit size. IRC spectroscopy is made at theslits located at the edge of the large imaging FoV. There isinternal scattered light in the array of the MIR-S and the light
diffusing from the imaging FoV affects SG1 and SG2 spectra
to some extent. This effect has been corrected for according to
the method given in Sakon et al. (2008). A similar scattered
light is also recognized in LG2 spectra and has been corrected
in a similar way. In addition, there is a contribution from thesecond-order light from the edges of the imaging FoV in LG2
spectra, but we expect that it is not significant in the backgroundsubtracted spectra.
These spectroscopic observations were simultaneously accompanied by supplemental short (49 s × 3) L18W-band imaging observations of a field ∼5 away from the central slit positionin the southeast direction in order to check the pointing of thesatellite. This supplemental, short-exposure L18W-band imageof the southeast, which fortuitously revealed interesting structures beyond the SNR boundary (see Section 3.4), was combinedwith deep (442 s) L18W-band image of the source to generate afinal larger-area map. The L18W-band imaging data were processed with the same procedures that we used for other imaging
data sets as described above.
3. RESULTS
3.1. Multi-band Infrared Images of G292.0+1.8
Figure 1 presents our AKARI IRC multi-band images ofG292.0+1.8 together with the Australia Telescope Compact
Array (ATCA) 20 cm radio continuum image (Gaensler &Wallace 2003), the Chandra X-ray (0.3–8 keV) image (Park443

et al. 2002), and the background-subtracted AKARI S11-band
image (S11–S7; see below) for comparison. The AKARI IR
emission associated with the SNR is most apparent in the L15and L24-band images as two prominent features: first, there isa ring-like structure composed of two clumpy, narrow, and long
filaments crossing the central part of the SNR along the east–west direction—we name this “equatorial ring” (ER); second,the outer boundary of the SNR appears as an almost-complete
shell structure with its eastern part opened—we name this “outerelliptical shell” (OES).
The radius of the ER is ∼3 (or ∼5 pc). Its southern filamentis brighter than the northern one by a factor of ∼1.7 on averageand has the brightest clump at (R.A., decl.) = (11h 24m 29.s 4,−59◦ 15 48 ) close to the filament center. The southern filamentalso shows prominent X-ray and [O iii] emission and has beencalled “equatorial bar” or “equatorial belt” (Park et al. 2002;
Gonzalez & Safi-Harb 2003; Ghavamian et al. 2005). The
northern filament is ∼1 away from the southern filament, andboth filaments are elongated roughly parallel to each other.
The OES shows an almost-complete ellipse of ∼7 × 6(∼12 pc × 10 pc) with its major axis aligning roughly withthe plane of the ER. The center of the ellipse determined byelliptical fits weighted by the 24 μm surface brightness of the
OES is (11h 24m 30.s 9, −59◦ 15 21 ) located at the middle of two
filaments of the ER. In the west, the OES appears to be connectedwith the two filaments of the ER, and its emission is enhanced
there, especially at the southwest near the southern filament of
the ER. In contrast, the eastern part of the OES is open, and also
the emission from the southern filament of the ER is truncatedin the east before it reaches the boundary of the SNR.
In addition to the ER and OES, there appear to be at least two
more clearly identifiable features in the AKARI images. First,
the L15-band image shows a separate structure that extends ∼5
southward from the ER on the southeastern side of the SNR. It iselongated vertically with the bright portion located ∼0. 5 below
the southern equatorial filament. Next, there is faint emission
extending beyond the northern and southeastern edges of the
OES. This faint emission is contained within the outermost
boundary of radio emission observed in G292.0+1.8 (Gaensler& Wallace 2003).
As in Figure 1, G292.0+1.8 is not detected in the S7- and S11band images: while some of the S11-band emission appears toarise from locations of strong L15- and L24-band emission,
there is almost no corresponding emission in the S7-band
image. This suggests that the S7-band emission is dominated bybackground emission, probably emission from PAHs in the lineof-sight interstellar medium (ISM). In the S11–S7 difference
image (Figure 1), we subtract out scaled S7-band emission fromthe S11-band image and confirmed that the two filaments of the
ER and southwestern part of the OES have appreciable S11band emission. (Note that the three bright point-like sources
close to the eastern end of the southern filament of the ER in the
S11–S7 difference image are stellar sources.) The NIR N3- and10 http://www.ir.isas.jaxa.jp/ASTRO-F/Observation/DataReduction/IRC/

Table 4
Slit Positions of AKARI Spectroscopic Observations
Slit Position

NG/SG1/SG2

LG2

Equatorial peak
Ejecta peak
Background(11h 24m 29.s 1, −59◦ 15 54 )(11h 24m 44.s 3, −59◦ 16 40 )(11h 25m 03.s 6, −59◦ 16 24 )(11h 24m 29.s 9, −59◦ 16 04 )(11h 24m 45.s 4, −59◦ 16 51 )(11h 25m 03.s 8, −59◦ 16 40 )Note. The slit position angles were fixed to be 44◦ by the satellite orbit.

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N3

N4

S7

S11

L15

L24

20cm

X-ray

S11-S7

Declination (J2000)

12

16-59° 20′

Declination (J2000)

12

16-59° 20′

Declination (J2000)

12

16-59° 20′11h 25m 00s
24m 30s
00s
Right Ascension (J2000)

11h 25m 00s
24m 30s
00s
Right Ascension (J2000)

11h 25m 00s
24m 30s
00s
Right Ascension (J2000)

Figure 1. AKARI IRC images of G292.0+1.8 with the ATCA 20 cm continuum and Chandra X-ray map for comparison (Gaensler & Wallace 2003; Park et al. 2002).
The background-subtracted S11-band image (S11–S7 difference image) is shown in the lower right panel. The three circles close to the eastern end of the ER in
the S11–S7 difference image indicate stellar sources. We mark the position of pulsar PSR J1124−5916 in N3-, L15-band, and 20 cm continuum images by the star
(Hughes et al. 2003). We also mark the center of the OES by the plus sign in the L15 image. The dotted box in the X-ray map indicates the imaging area of Chandra.
The gray scales of N3, N4, S7, S11, L15, L24, and S11–S7 bands are 0.5–3.0, 0.5–3.0, 5.3–6.3, 15.8–16.7, 23–26, 29–40, and 0–0.4 MJy sr−1 , respectively.

N4-band images, on the other hand, show only stellar emission
without any apparent feature brighter than 16 μJy at 3 μm thatmight be associated with the SNR.
Table 5 presents IR measurements of G292.0+1.8. In the table,
we list the fluxes and ratios of whole remnant area and IR-Ejecta
region (see befow for IR-Ejecta). In addition, we also list thepeak intensities and ratios of the ER, the southwestern OES, andthe IR-Ejecta.
3.2. Infrared Colors and Ejecta Identification
Figure 2 (left), which is a three-color (7, 15 and 24 μm)
AKARI MIR image of G292.0+1.8, shows that most of theprominent features have roughly similar MIR colors. One
exception is the feature in the southeastern part of the SNR
with significant excess in the shorter wavebands—we name thisfeature “IR-Ejecta” because it is believed to originate in shocked
SN ejecta as we describe below. It is bright near the southernequatorial filament and stretching directly southward beyond
the SNR boundary. This feature is clearly seen in Figure 2(right) which shows the L15/L24 emission ratio. The ratio
image was produced by dividing the L15-band image by theL24-band image after background subtraction. The backgroundwas estimated by fitting a slanted plane to the areas surrounding
the SNR. The resulting background planes were almost flat, i.e.,the brightness differences over the entire images were only 1.7%
and 2.4% in L15 and L24, respectively. We applied a Gaussianconvolution to the L15-band image in order to match the final
spatial resolution to that of the L24-band image and masked out
the pixels with small (30.2 MJy sr−1 ) L24-band intensities.
The resulting L15/L24 color image is significantly different
from the original band images. In Figure 2, the L15/L24 ratio isalmost uniform (∼0.25), and there is no feature corresponding
to the prominent ER and OES. The most notable feature is
the IR-Ejecta, where L15/L24 ratio raises to ∼0.8, which is infact consistent with what we identified in the three-color image

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445

Table 5
IR Measurements of G292.0+1.8
Band

Whole Remnant

Measured flux/intensity
S11
L15
L24
N60
Wide-S
Wide-LColor
S11/L15
L15/L24
N60/Wide-S
Wide-S/Wide-L

IR-Ejecta

Total Flux
(Jy)

ER Peak Intensity(MJy sr−1 )OES Peak Intensity(MJy sr−1 )Ejecta Fluxa(Jy)

Peak Intensityb(MJy sr−1 )0.30 (0.02)
2.81 (0.14)
11.98 (0.72)
30.6 (6.1)
30.0 (9.0)< 6.8

0.53 (0.26)
5.10 (0.36)
19.09 (1.37)
8.81 (2.53)
6.76 (2.14)< 6.500.53 (0.24)
5.04 (0.36)
18.94 (1.26)
7.65 (1.83)
6.97 (2.20)< 6.580.013 (0.001)
0.066 (0.009)
0.072 (0.015)
0.173 (0.035)
0.250 (0.075)< 0.048

0.37 (0.24)
1.96 (0.26)
2.60 (0.36)
5.32 (1.47)
6.05 (1.95)< 6.87

0.11 (0.01)
0.23 (0.02)
1.02 (0.37)> 4.420.10 (0.05)
0.27 (0.03)
1.30 (0.56)> 1.040.11 (0.05)
0.27 (0.03)
1.10 (0.43)> 1.060.20 (0.03)
0.92 (0.23)
0.69 (0.25)> 5.250.19 (0.13)
0.75 (0.14)
0.88 (0.37)> 0.88Notes.
a The area used for the flux measurement is indicated in the inset of Figure 3.
b Measured intensity at the peak of SG slit position.
0.0

Declination (J2000)

12

14

0.2

0.5

0.7

OES
ER

16

18

-59° 20′102

11h 25m 00s

24m 30s
Right Ascension (J2000)00s

IR-Ejecta

11h 25m 00s

24m 30s
Right Ascension (J2000)

00s

Figure 2. Left: three-color image of G292.0+1.8 produced from 24 μm (R), 15 μm (G), and 7 μm (B). The coordinates of equatorial, ejecta, and background
NG/SG slit positions are indicated by the white box at (11h 24m 29.s 1, −59◦ 15 54 ), (11h 24m 44.s 3, −59◦ 16 40 ), and (11h 25m 03.s 6, −59◦ 16 24 ), respectively. At theirsouthwestern side, the three LG2 slit positions are indicated by a gray box at (11h 24m 29.s 9, −59◦ 16 04 ), (11h 24m 45.s 4, −59◦ 16 51 ), and (11h 25m 03.s 8, −59◦ 16 40 ),respectively. We mark the center of the OES by plus sign, both the dynamical center of optical ejecta knots (Winkler et al. 2009) and the center of outer radio plateau(Gaensler & Wallace 2003) by crosses, and the position of pulsar PSR J1124−5916 by a star (Hughes et al. 2003). The two cross symbols are almost coincident.
Right: the background-subtracted 15/24 μm ratio image. The pixels with low 24 μm intensities are blanked out. The bright structures in 24 μm images are shown by
contour. The prominent features such as the ER, OES, and IR-Ejecta are also labeled.

(the left panel of Figure 2). It appears to be of a triangularshape with one of its vertices toward the direction to the SNR
center. The southern vertex passes though the southern boundaryof the OES and extends over the SNR boundary, while thenortheastern vertex is located just above the ER. The region ofhigh color ratio of the IR-Ejecta covers a large portion of thesoutheastern area of the SNR and positionally coincides with theO- and Ne-dominant ejecta identified by previous optical and
X-ray observations (Ghavamian et al. 2005; Winkler & Long
2006; Park et al. 2002). There is also a wispy patch of emission
around (11h 24m 43.s 4, −59◦ 20 47 ) far (> 1 ) beyond the SNR

boundary, which has the color ratio similar to that of the IREjecta. It, however, is not shown in the [O iii] images (Winkler& Long 2006). Besides the IR-Ejecta in the southeastern region,
the northwestern region shows extended emission of the highcolor ratio, although it is not as significant as the southeasternregion. Overall the two regions of the high color ratio are roughly
symmetric with respect to the center of the OES.
Figure 3 compares the pixel values of the L24- and L15-band
emission, where most of the points are distributed along the
thick, solid line of a slope of 0.25. (Note that pixels of stars
and pixels with small L24-band intensities are masked out.) If

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Figure 3. L15 vs. L24 band intensity plot. We mask the pixels with stellar emission using the N3 band and the pixels with low intensity using the L24 band. The
vertical lines indicate the boundaries of the high L15/L24 ratio points with tight correlation (emission only from ejecta). The thick line is a linear fit using all the
pixels in the plot. The thin line is a linear fit using the high L15/L24 ratio region between two vertical lines. The position of high L15/L24 ratio pixels are marked in
the inset in the lower right corner superposed on the L24 contour. Black dots are the positions of the high L15/L24 ratio pixels between two vertical lines, and graydots are those of high L15/L24 ratio pixels with L24 intensity  33 MJy sr−1 (ejecta emission superposed on swept-up material).

L24- and L15-band fluxes are from the same region and if they
are both thermal dust emission, then the color ratio and the dust
temperatures are related asIν (15)κν (15)Bν (15, T )=,Iν (24)κν (24)Bν (24, T )(1)where Iν (λ) is the surface brightness at λ (μm) in Jy sr−1 , Bν (T )is the Planck function, and κν (λ) is the dust opacity in cm2 g−1 .
The slope of 0.25 corresponds to a dust temperature of 126 K fora mixture of carbonaceous and silicate interstellar grain of RV =
3.1 (Draine 2003). In Figure 3, there are two regions, where thedata points show a high L15/L24 ratio compared to the thick,solid line: first, the thin line of the slope of 0.88 represents
one group whose L24-band surface brightness is less than
∼33 MJy sr−1 ; second, there is another group of data points ofhigher L15-band emission whose L24-band surface brightness
lies in the range of 33–40 MJy sr−1 . The inset in the lower-rightcorner of Figure 3 shows that the data points with high L15/L24
ratios are all from the IR-Ejecta as expected. The ones in thesecond group, i.e., ones with high 24 μm surface brightness,
are superposed on the ER, and therefore their emission is partly
from the swept-up CSM while the data points in the first groupshould represent the emission only from the ejecta. The slope of
0.88 corresponds to a dust color temperature of 240 K; however,the L15-band flux is largely from line emission at this position
so that the physical dust temperature should be lower than this(see Section 3.4).

Figure 4. Background-subtracted spectra obtained at the equatorial peak andejecta positions. The MIR wavelength is covered by three elements: SG1 (4.6–9.2 μm), SG2 (7.2–13.4 μm), and LG2 (17.5–26.5 μm). Note that the positionsof LG2 spectra are slightly shifted from those of SG spectra. Their positions
and coordinates are indicated in Figure 2. The relative spectral responses of S7,
S11, L15, and L24 imaging bands are indicated in the bottom panel.

3.3. Mid-infrared Spectroscopy of Ejecta and Equatorial PeakFigure 4 shows the background-subtracted MIR grism spectraof the peak positions of the ER and the IR-Ejecta.11 (Note that the
11 We only present the MIR spectra because NIR spectra were heavily

contaminated by emission from nearby field stars.

LG2 spectra at >17 μm were obtained from positions slightly
shifted from those of the SG spectra of 5–14 μm as we havedescribed in Section 2.) The spectrum from the IR-Ejecta peak
clearly shows the [Ne ii] line at 12.8 μm, which confirms thenature of the radiatively shocked ejecta. No Ar lines (i.e., [Ar ii]

No. 1, 2009AKARI INFRARED OBSERVATIONS OF THE SNR G292.0+1.8

447

Table 6
Results of AKARI IRC Spectroscopy
Line

Wavelength(μm)

ER Peak(erg cm−2 s−1 sr−1 )IR-Ejecta Peak(erg cm−2 s−1 sr−1 )[Ar ii][Ar iii][Ne ii]

7.0
9.0
12.8

< 1.6 × 10−6
< 4.0 × 10−7
< 2.7 × 10−6

< 2.9 × 10−6
< 1.1 × 10−6
2.8 × 10−5Note. Undetected lines below 3σ detection limit.

at 7.0 μm and [Ar iii] at 9.0 μm) were detected, while neither[Ne iii] line at 15.6 μm nor [O iv] line at 25.9 μm was covered.
Table 6 lists the flux of the [Ne ii] line and upper limits of
the [Ar ii] and [Ar iii] lines. The flux of the [Ne ii] line at
12.8 μm is 2.8 ± 2.5 × 10−5 erg cm−2 s−1 sr−1 , corresponding
to the surface brightness of 0.24 ± 0.21 MJy sr−1 in the L15band image. Given the L15-band surface brightness of 1.6 ±0.1 MJy sr−1 , ∼15% of the L15-band emission is due to the[Ne ii] line emission. For the S11 band, the observed [Ne ii] line
emission flux corresponds to a S11-band surface brightness of
0.12 ± 0.11 MJy sr−1 . This is equivalent to ∼55% contribution
to the total S11-band emission, which is much larger than thecase of the L15-band emission. The difference is because there
is no strong emission line other than the [Ne ii] line in the S11band while there is an additional [Ne iii] 15.5 μm line in theL15 band (see the spectral responses in Figure 4). (There could
be some dust emission too; see Section 4.3.) On the other hand,the spectrum at the ER peak is dominated by the continuum
emission, although the LG2 spectra have a lower signal-to-noise
ratio.
3.4. Large-scale Infrared Emission Around G292.0+1.8
Figure 5 presents a combined L18W-band image (Section 2)
covering both the SNR and an extended area in the southeast.
Besides the features of the SNR that we have already describedin previous sections, there is a notable feature of the L18-band
emission in the south. The elongated “narrow tail” is locatedaround (11h 24m 45s , −59◦ 22 00 ), ∼7 apart from the center of
the SNR. It is close to the IR-Ejecta in the southeastern partof the SNR and its elongation direction is roughly toward thecenter of the SNR. Its extent is ∼1. 5 or 2.6 pc. In addition,
there is faint, diffuse emission toward the south and southeast
too. The emission toward the southeast appears distinct—it
appears to protrude from the open portion of the OES having
a similar outer boundary with radio plateau at the east andsouth, but extends ∼7 southeast beyond the radio boundary of
the SNR. We consider that this “wide tail” could be associatedwith the SNR. There is a large, diffuse H ii region superposedon G292.0+1.8 on the sky (RCW 54; Rodgers et al. 1960),
but its Hα emission is elongated along the northeast–southwest
direction without any emission corresponding to the wide tailin the Southern H-Alpha Sky Survey Atlas (SHASSA; Gaustad
et al. 2001). Therefore, the wide tail is not associated with H ii
regions and its association with the SNR is likely.
Figure 6 presents the AKARI FIS FIR-band images ofG292.0+1.8. The images cover the entire SNR with a scan direction from northeast to southwest. The FIR emission of the
SNR associated with the ER and OES is clearly detected in
the N60- and Wide-S-band images with strong concentrationin the southwestern part, although fine details are difficult tosee because of their low spatial resolutions. The N160- andWide-L-band images are, however, clearly different from images of the other bands. They show an elongated feature ex-

Figure 5. Large-scale L18W-band image of G292.0+1.8. The gray-scale range
is indicated at the top. Contour levels are 27.0, 27.5, 29, and 31 MJy sr−1 .
The dotted contour is the ATCA 20 cm radio-continuum boundary (Gaensler& Wallace 2003). The “narrow tail” and “wide tail” features outside the shellboundary are labeled.

tended in the northwest–southeast direction which shows nocorrelation with the emission associated with the SNR. Also
its peak position is located outside the boundary of the SNR.
These indicate that the N160- and Wide-L-band emission is
not directly associated with G292.0+1.8, although it is possiblethat the enhanced brightness in the SW part of the OES is dueto the interaction of the remnant with this extended structure(cf. Braun et al. 1986). The bottom panels of Figure 6 present
the background-subtracted N60- and Wide-S-band images. The
background emission was estimated by calculating scaling factors between the N60- and N160-band images and also between
the Wide-S- and N160-band images from the areas outside the
SNR. Compared to the N60- and Wide-S-band images of the top
panels, the background-subtracted images show more clearly theemission associated with the SNR, including the northern partof the OES that is not clear in the original images.
4. DISCUSSIONS
4.1. Destruction of Circumstellar Dust
4.1.1. Infrared Emission from Shocked Circumstellar Dust

The ER and OES are the most prominent features in our
AKARI multi-band IR images (Figures 1 and 6). The ER is
composed of two filaments, where the southern one is brighter
than the northern counterpart. The southern filament appears to
be of the normal composition in soft X-rays (Park et al. 2002)
without any apparent radial motion in the optical (Ghavamian
et al. 2005), which led the authors to conclude that it is CSM
from the progenitor of the SN in G292.0+1.8. We suggest thatthe northern filament, which has clearly been found by AKARI
observations, could be part of the same structure based on
its similar distribution to the southern filament. We note thatthe bright clump at (11h 24m 46s , −59◦ 14 42 ) near the easternend of the northern filament has a counterpart in the [O iii]
image of Ghavamian et al. (2005). Its velocity is near zero,LEE ET AL.

Declination (J2000)

448

Vol. 706

N60

Wide-S

Wide-L

N160

N60 - N160

Wide-S - N160

15

Declination (J2000)-59° 25′

15

Declination (J2000)-59° 25′

15-59° 25′
11h 25m
24
Right Ascension (J2000)

11h 25m
24
Right Ascension (J2000)

Figure 6. AKARI FIS images of G292.0+1.8. The scan direction is northeast to southwest. The observed band is indicated at the upper right corner in each panel.
The background-subtracted images are shown in the bottom panels. The gray scales of N60, Wide-S, Wide-L, N160, N60–N160, and Wide-S–N160 bands are 18–28,
32–44, 70–88, 52–64, 0–9, and 0–8 MJy sr−1 , respectively.

which supports that the northern and southern filaments forma single structure. Furthermore, it is most clearly visible at the−120 km s−1 frame of the Rutgers Fabry–Perot velocity scan
images and absent at the 0 km s−1 frame where the southern
filament is brightest (Figure 2 of Ghavamian et al. 2005). We
interpret this velocity difference as an indication suggestingthat the northern and southern filaments are parts of a tilted,
expanding ring structure produced by the progenitor of the SN.
The location of the center point of the OES at the middle of thetwo filaments also reconciles with the interpretation, given the
CSM nature of the OES (see below). We note that such a, but
smaller, ring structure of the CSM was found in SN 1987A andalso possibly in the Crab nebula (Bouchet et al. 2004; Greenet al. 2004). The OES, on the other hand, is relatively fainter in
X-rays than the ER (Park et al. 2002). The AKARI MIR
brightness and color of the OES, however, are similar to thoseof the ER (Figure 2). Toward the southern filament of the ER,
Ghavamian et al. (2005) reported the detection of the opticalradiative lines produced by the partially radiative shocks starting
to develop cooling zones. It implies that the X-ray emitting
gas in the ER is cooler than that of the OES, while they havesomewhat similar dust properties. And there is a possibility that
the ER in the S11–S7 difference image also contains the [Ne ii]
line emission produced in the CSM region where the shock has
started to cool down to <100,000 K. However, its contribution
might be small, because we have not detected the [Ne ii] 12.8 μm
line at the ER.
In addition to the ER and OES, there are faint MIR emissions beyond the northern and southeastern edges of the OES(Figure 2). They have sharp outer boundaries, representing thecurrent locations of the SN blast wave. Faint X-ray emission
was detected in those areas too (Park et al. 2002). It is possiblethat the SN blast wave has overtaken the OES and produced
those emission features while propagating into a more diffuse
medium. On the other hand, it is also possible that the remnant has a front–back asymmetry and they are just projected

No. 1, 2009AKARI INFRARED OBSERVATIONS OF THE SNR G292.0+1.8

Figure 7. Measured total fluxes and the best fit of SED using the modified
blackbody composed of two temperature dust components. The dust model isa mixture of carbonaceous and silicate interstellar grains of RV = 3.1 (Draine
2003). The boxes are the measured fluxes with errors, and the triangles indicate
the color-corrected fluxes. The downward arrow indicates the upper limit of the
140 μm flux.

boundaries of the more-extended shell. In any case, their asymmetric spatial distribution suggests that either the ambient density distribution or the mass ejection from the progenitor was
asymmetric.
As in Figure 4, the MIR emission of G292.0+1.8 is dominatedby continuum emission, not by line emission. This implies that
the IR emission is from shock-heated dust grains in the CSM.
Figure 7 presents the spectral energy distribution (SED) of the
SNR in Table 5. We applied modified blackbody fits composedof two dust components to the observed SED in order to obtain
the best-fit dust temperatures. The dust model based on a mixtureof carbonaceous and silicate interstellar grain of RV = 3.1(Draine 2003) gave dust temperatures of ∼103 K (warm dust)and  47 K (cold dust), corresponding to the mass of 4.5 ±0.9 × 10−4 M and  4.8 × 10−2 M , respectively. (Notethat the lower limit of the cold dust temperature comes from theupper limit of the flux at 140 μm.) For the case of the dust model
based on the graphite and silicate grain of 0.001–0.1 μm size
(Draine & Lee 1984; Laor & Draine 1993), the derived total
mass is in the range of (1.0–3.4) × 10−2 M , comparable tothe total mass obtained for the former model. The derived dust
mass corresponds to a dust-to-gas ratio of 1.6 × 10−3 , if we
use the 30.5 M swept-up mass (at a distance of 6.2 kpc) of
the gas in the CSM obtained in X-ray observations (Gonzalez &Safi-Harb 2003). This is lower than the ratio 6.2 × 10−3 foundin the local ISM (Zubko et al. 2004). The low dust-to-gas ratios
in the swept-up materials by the shock destruction were also
obtained with Spitzer observations on the core-collapse SNRs
in the Large Magellanic Cloud (LMC; Williams et al. 2006) andthe Kepler (Blair et al. 2007). According to our result, 75%of the dust in G292.0+1.8 might have been destroyed by the SN
shock, or the initial dust-to-gas ratio surrounding G292.0+1.8
might be lower than the local value. It is comparable to thefraction derived in other SNRs, such as 64% in Cas A (Dweket al. 1987) and 78% in Kepler (Blair et al. 2007).
4.1.2. Model Calculations of Shock-heated Dust Emission

We perform model simulations for the destruction of dust by
SN shock waves and the thermal emission from shock-processed
dust. The physical processes of dust in shocks have so far been
discussed in many works (e.g., Tielens et al. 1994; Vancura et al.
1994; Dwek et al. 1996; Jones 2004). Once the circumstellar

449

dust grains are swept up by the blast wave, they acquire highvelocities relative to the gas and are eroded by kinetic and/or
thermal sputtering in the shock-heated gas. Dust grains are also
heated by collisions with energetic electrons in the hot gas andradiate thermal emission at IR wavelengths. Dynamics, erosion,and temperature of dust depend on the temperature and densityof the gas as well as the chemical composition and size of dust
grains.
We adopt the model of dust destruction calculation by Nozawa
et al. (2006), in which the motion, destruction, and heating of
dust in the shocked gas are treated in a self-consistent manner byfollowing the time evolution of the temperature and density of
the gas for the spherically symmetric shock wave. As the initialcondition of the SN ejecta we consider the freely expanding
ejecta with a velocity profile of V = Veje (r/Reje ) and a density
profile of ρ = ρcore at r  Rcore and ρ = ρcore (r/Reje )−α atr > Rcore , where Veje and Reje are the velocity and radius of
the outermost ejecta, respectively. Taking the kinetic energy of1051 erg and the ejecta mass of 8 M (Gaensler & Wallace 2003)
for the density profile of core-collapse SNR of Rcore = 0.3Rejeand α = 12 (e.g., Chevalier 1982; Matzner & McKee 1999;
Pittard et al. 2001), we obtain ρcore = 1.4 × 10−19 g cm−3 ,Veje = 1.3 × 109 cm s−1 , and Reje = 4.2 × 1017 cm at 10 years
after explosion, when the simulations are started. Since most of
the dust grains swept up by the forward shock during the later
epoch of the evolution, the calculation results are not sensitive
to the ejecta structure.
For the ambient medium we consider the constant hydrogen
number density of nH,0 = 0.1, 0.5, 1, and 10 cm−3 . The
circumstellar dust is assumed to be amorphous carbon or silicate(forsterite) with the power-law size distributions (∝ a −3.5 )ranging from a = 0.001 μm to a = 0.5 μm (Mathis et al.
1977). The optical constants are taken from Edo (1983) and
Semenov et al. (2003). Based on the time evolution of the size
distribution and temperature of dust given by the simulation andthe assumption of the dust grains being in thermal equilibrium,
we calculate the IR SED by thermal emission from the shockheated dust. The detailed description for calculating the IR SED
from shocked dust will be given elsewhere (Nozawa et al. 2009).
Figure 8 compares the observed fluxes of G292.0+1.8 withthe calculated IR SEDs at 3000 years for nH,0 = 0.1, 0.5, 1, and10 cm−3 . We present the results with the initial dust-to-gas mass
ratio to best reproduce parts of the observed SED, for amorphous
carbon (Figure 8(a)) and silicate (Figure 8(b)). The typicaltemperatures of dust are 35–55, 45–65, 50–70, and 60–80 K for
nH,0 = 0.1, 0.5, 1, and 10 cm−3 , respectively, and the resulting
dust masses are in the range of (0.3–5) ×10−2 M for carbon and(0.4–8) ×10−2 M for silicate with the higher values for lower
ambient density (thus lower temperature of dust). It can be seenthat the results of the silicate grains with nH,0 = 0.5 cm−3 , whichcoincide with the density estimated from X-ray observations(Gonzalez & Safi-Harb 2003), can reasonably reproduce theoverall shape of the IR SED for G292.0+1.8. In this case,the mass of grains radiating IR emission is 2.9 × 10−2 M .
Note that the initial dust-to-gas mass ratio corresponding thisresult is 1 × 10−3 , which is smaller than that in the local
Galaxy. If the IR emission originates in the swept-up materials
containing the mass-loss wind of red supergiant (RSG) with thesolar metallicity, the condensation of silicate is expected, and,according to our results, its condensation efficiency could be
low.
It should be noted here that the simulation results for nH,0 <10 cm−3 significantly underestimate the flux at short (11 μm)

450

LEE ET AL.

Vol. 706

Figure 8. Comparison of the measured fluxes with the SED at 3000 years obtained from the model calculations (see the text for details.). The dust compositions are(a) amorphous carbon grains and (b) silicate grains. The dotted, solid, dashed-dotted, and dashed lines indicate the results of simulations for nH,0 = 0.1, 0.5, 1, and10 cm−3 . Symbols are the same as Figure 7.

wavelength. However, the SED at shorter wavelengths could be
resolved by including the effect of a stochastic heating of small
grains; stochastically heated dust produces more emission atshorter wavelengths than the dust with equilibrium temperature,and it may also allow acceptable fits for lower initial density and/or different size distributions of dust than the current best fit.
Alternatively, this disagreement may be caused by the differencein the assumed composition of dust. To gain deeper insight
into the properties of dust in G292.0+1.8, we need further
investigations by taking account of the stochastic heating andchanging the composition and size distribution of dust as wellas the density profile in the ambient medium.
4.2. Infrared Emission from Supernova Ejecta
4.2.1. Ejecta Neon Line Emission

We have detected MIR emission associated with the SN
ejecta: IR-Ejecta. It is prominent in the 15/24 μm ratio image
by its high L15/L24 ratio (Figure 2), but is marginally seen in
the total intensity maps of the MIR (11–24 μm) and FIR (65–
90 μm) too. The emitting area coincides with the fast-moving(<1500 km s−1 ), O-rich SN ejecta: the triangle-shaped bright
portion near the equatorial filament was called “spur” and theextension to the south was called “streamers” by Ghavamian
et al. (2005). In the high-resolution [O iii] 5007 image, thespur is crescent shaped with a sharp boundary toward the SNR
center, while the streamer appears to be composed of clumps
embedded in diffuse emission. Recent measurement of theirproper motions showed that they are expanding systematicallyfrom a point near the geometrical center of the OES (Winkler
et al. 2009, see Section 4.4 as well).
According to our spectroscopic result, the IR-Ejecta shows[Ne ii] 12.8 μm emission but no [Ar ii] 7.0 or [Ar iii] 9.0 μm
emission. Note that the latter lines from Ar ions were also
detected at the metal-rich ejecta in young SNRs (Douvion et al.
2001; Smith et al. 2009; Williams et al. 2008). The lack of IR line
emission from Ar ions in this area is consistent with the resultsof optical or X-ray studies which showed that there is no line
emission from elements heavier than S in this area (Ghavamian
et al. 2005 and references therein; Park et al. 2002). The absenceof Ar lines supports the claim by Ghavamian et al. (2005) that
we are not seeing the ejecta in the innermost region accelerated
by PWN but seeing the He-burning-synthesized, O-rich ejecta
swept-up by the reverse shock.

The IR-Ejecta shows high L15/L24 ratio, i.e., L15/L24 =0.88 compared to 0.25 of the shocked CSM. The high L15/
L24 ratio is at least partly due to Ne lines. (There is a [O iv]
25.9 μm line in the L24 band, but it is weak, i.e., ∼20% of
the [Ne ii]+[Ne iii] lines; see Section 4.3.) According to ourestimation, the [Ne ii] 12.8 μm line contributes ∼15% to theL15 flux. In the L15 band, there is another strong Ne line, the[Ne iii] 15.6 μm line. We may estimate the possible contributionof the [Ne iii] 15.6 μm line in G292.0+1.8 as follows.
The strength of a forbidden Ne line is given by I21 =(hν21 /4π )N2 A21 , where ν21 is the frequency of the line, N2
is the column density of Ne+ or Ne++ ions in the upper statealong the line of sight, and A21 is the Einstein coefficient, so thatthe [Ne iii] 15.6 μm/[Ne ii] 12.8 μm ratio is given by
I15.6 μm
N (Ne++ ) f (Ne++ , 3 P1 )= 0.57,I12.8 μm
N (Ne+ ) f (Ne+ , 2 P1/2 )(2)

where N (Ne++ ) and N (Ne+ ) are column densities of Ne++ andNe+ ions, f (Ne++ , 3 P1 ) and f (Ne+ , 2 P1/2 )are the fractions ofNe++ and Ne+ ions in the upper states of the corresponding
emission lines (see Glassgold et al. 2007 for atomic parametersof these lines). Assuming statistical equilibrium, the last factor,f (Ne++ , 3 P1 )/f (Ne+ , 2 P1/2 ), can easily be calculated (e.g.,
Glassgold et al. 2007), and varies from 1 to 2 as the electron
density increases from a low-density limit to densities muchhigher than the critical density (= (4–6) × 105 cm−3 at
10,000 K). Therefore, the intensity ratio I15.6 μm /I12.8 μm is  1
unless N (Ne++ )  N (Ne+ ), which happens at temperatures
105 K in collisional equilibrium (e.g., Allen & Dupree1969). Ghavamian et al. (2005) suggested that ejecta undergoes
radiative shocks with a velocity of 50–200 km s−1 from the[O iii] 5007 line widths. The temperature of shocked ejecta gas
immediately behind the shock of velocity v s should be high, i.e.,
4.5×106 (vs /100 km s−1 )2 K, assuming pure neon preshock gas,
but since the shocked gas element cools fast due to enhanced
heavy elements and the column density of the gas at T mightvary roughly with 1/T , we expect N (Ne++ ) < N (Ne+ ) andtherefore I15.6 μm /I12.8 μm  1. Toward the SN ejecta in young
core-collapse SNRs such as Crab, Cas A, the SNR B0540−69.3
in the LMC, or the SNR 1E0102−72.3 in the Small Magellanic
Cloud (SMC), the [Ne iii]15.6 μm/[Ne ii]12.8 μm ratio varies
from 0.3 to ∼1 too (Douvion et al. 2001; Temin et al. 2006;
Smith et al. 2009; Williams et al. 2008; Rho et al. 2009). We

No. 1, 2009AKARI INFRARED OBSERVATIONS OF THE SNR G292.0+1.8

may conclude that the contribution of the [Ne iii]15.6 μm line
to the L15 band should be at most comparable to that of the[Ne ii]12.8 μm line unless the shock is truncated at relatively
high temperatures.

451

We have derived IR fluxes associated with the ejecta andthe results are summarized in Table 5. It is difficult to extractthe ejecta emission in IR observations, although our results
unveiled the ejecta region unambiguously. The IR-Ejecta isclearly divided into two regions (Figures 2 and 3); one thatcoincides with the ejecta-only region and has a tight correlationin the L15 versus L24 plot, and the other superposed on the ER
and has a loose correlation in the L15 versus L24 plot. These
IR-Ejecta regions are part of the structure defined as the [O iii]
“spur” by Ghavamian et al. (2005), with the ejecta-only region
being the bottom of the spur, and the ejecta superimposed on
the ER being the top portion of the spur. We measure the flux
at the ejecta-only region masking the emission from stars. The
area of the ejecta-only region is more than four times larger than
the superposed ejecta region and it contains most of the flux.
In order to remove the background emission, we subtract the
average of the nearby (∼2 ) low-intensity region.
The AKARI ejecta spectrum in Figure 4 shows that the 11 and15 μm fluxes are not from dust continuum. We believe that theS11 band is mostly dominated by line emission, while the L15
band needs an additional component (see Section 3.3). In thefollowing section, we show that some of the L15 flux is due to abump at 15–25 μm. On the other hand, FIR emission is certainly
dominated by dust continuum. The L24 band contains the [O iv]
25.9 μm line emission, whose wavelength is not covered with
our spectroscopic observations. Its contribution in the L24 bandis expected to be less than 20% in our estimation using the[O iv] 25.9 μm line intensity in Ghavamian et al. (2009). At
first, we fit 24–140 μm band fluxes using a modified blackbody
for silicate and graphite grain models of 0.001–0.1 μm size
(Draine & Lee 1984; Laor & Draine 1993). The best fits give atemperature of 64–88 K and a dust mass of (2.0–8.2) × 10−5 M .
If we exclude the L24-band data point, the temperature drops
to 44–52 K and the dust mass increases by a factor of 5. This
IR-emitting dust mass is much smaller (< 4 × 10−4 ) than thetheoretically predicted dust mass (0.1–1 M ) to be formed in
the core-collapse SN explosion (e.g., Nozawa et al. 2003).

fluxes are entirely due to lines, we estimate that the L15/L24
ratio should be 2.3–2.6 using the line fluxes in Ghavamian et al.
(2009). (It becomes higher if we consider the 15–25 μm bumpin the following paragraph.) The observed ratio, however, is0.25 for the ER and OES and 0.88 for the ejecta. Therefore, thecontribution of the [O iv] 25.9 μm line flux to the observed L24band flux should be small, particularly toward the ER whereIν (24)/Iν (15) ≈ 5 (Section 3.2). We, however, note that theSpitzer IRS spectrum toward the ejecta in Ghavamian et al.
(2009) does not show any obvious dust continuum emission
between 24 and 36 μm. We consider that it could be because
there was dust continuum emission in the IRS background. In
the AKARI 24 μm image, there is faint filament coincident with
the IRS background region (LL1 Sky in Ghavamian et al. 2009).
This filament is associated with the OES and its 24 μm emission
might be dominated by dust continuum, so that, by subtracting
its spectrum from the ejecta spectrum, the continuum feature
could have been removed.
An interesting feature in the Spitzer ejecta spectrum is a weak
15–25 μm bump, which was suggested to be produced by newly
formed dust or swept-up PAHs along the line of sight. First ofall, this bump explains the high L15/L24 ratio at the ejecta
region described in Section 4.2.1. The peak of the bump featureappears in the L15 band. Therefore, it contributes mainly to the
AKARI 15 μm band. Second, our AKARI observations show thatthe areas with a high L15/L24 ratio are coincident with those of
the optical O-rich ejecta region. This supports the interpretationthat the bump feature is related to the newly formed dust in
association with the SN ejecta, not to the swept-up PAHs.
According to Ghavamian et al. (2009), the Spitzer spectrum
from the southern filament of the ER is consistent with theemission from two dust components: a warm (or hot) componentof ≈ 114 K and a cold component of ≈ 35 K. The temperatureof the warm dust is higher than what we estimated with the
AKARI data by ∼10 K, while that of the cold dust is lower thanthe AKARI temperature by a comparable amount. It is possiblethat the presence of the bump feature in the Spitzer spectrum
resulted in a higher temperature for the warm dust. For the cold
dust temperature, as explained in Ghavamian et al. (2009), theabsence of the longer wavelength (>30 μm) data in the Spitzer
spectrum is likely the primary reason of the lower temperaturecompared to the AKARI results. IR spectroscopic observations
covering a broad band are necessary to clearly resolve the issues.

4.3. Comparison with Spitzer Spectroscopic Results

4.4. SN Explosion in G292.0+1.8

Recently, the results from MIR observations of G292.0+1.8
with the Spitzer satellite have been published (Ghavamian et al.
2009) based on low-resolution spectroscopy of the ejecta andthe southern filament of the ER where the ejecta emission
is superimposed on the emission from the CSM. The former
is located ∼0. 5 north from our IR-Ejecta slit position, while
the latter, which has a high L15/L24 ratio, is ∼1 apart in
the northeast direction. Their ejecta spectrum shows strong
emission from the [Ne ii] 12.8 μm, [Ne iii] 15.6 μm, [O iv]
25.9 μm lines and relatively weak [Ne v] 24.4 μm and [Ne iii]
36.0 μm lines, but no lines of Mg, Si, S, Ar or Fe are identified.
(Note that the Si and S lines in their spectrum are not from theejecta but from the background.) The observed ratio of [Ne ii]
12.8 μm to [Ne iii] 15.6 μm is 2.1–2.7. The non-detection of
Ar lines and the observed ratio of Ne lines are consistent with
our results and the prediction in Section 4.2. The observed fluxof the [O iv] 25.9 μm line is ∼20% of the [Ne ii] plus [Ne iii]
line fluxes (Ghavamian et al. 2009). If the AKARI L15 and L24

4.4.1. Circumstellar Shell and the Explosion Location

4.2.2. Dust Emission Associated with Ejecta

One interesting result that we obtained in this study is thedifference (∼42
1 pc) between the center of the CSM
and the dynamical center of the ejecta (Figure 2). The former
corresponds to the center of the ER and OES that we determinedwith the AKARI results; the latter was determined by the
distribution of the O-rich ejecta in the optical, which is close
to the center of the SNR in the radio emission (Winkler et al.
2009). In addition, the position of the pulsar is different fromthe both positions—it is shifted in the southeast direction
from the dynamical center by ∼46 .
One possibility is that the progenitor star was at the centerof the OES during its RSG phase but exploded at the positionof the dynamical center of optical knots. This is possible
because the RSG wind could be confined by external pressure
while the central star is moving. According to Chevalier (2005),
the RSG wind from a 25–35 M star, which explodes as

452

LEE ET AL.

SN IIL/b, would be pressure confined while its outer radius
becomes  5 pc. The radius of the OES (6 pc) is comparable
to what the theory predicts. If the progenitor star was moving
at ∼10 km s−1 and exploded after ∼105 yr of the pressure
confinement of the shell, then the explosion center would beclose to the dynamical center of optical knots.
Another possibility is that the SN exploded at the center ofthe ER and OES, not at the dynamical center of the ejecta or thecenter of the radio emission. This is motivated by the fact that
the OES shows a very well-defined shell structure surrounding
the SNR, thereby the center of the OES may pinpoint to thereal location of the progenitor, which later exploded as a SN.
On the other hand, the centers of the radio emission and the
O-rich ejecta motion could have been weighted toward thesoutheast: firstly, the geometrical center of the radio nebula
could be significantly weighted toward the southeast becauseof the existence of the bright PWN. Second, the dynamical
center of optical knots is also weighted to the southeast because
the bright optical knots in the southeastern area are moving
relatively slowly and the center is derived using an assumptionof an unhindered constant velocity since the explosion (Winkler
et al. 2009).
If the SN explosion in G292.0+1.8 indeed occurred at thecenter of the OES, then the tangential velocity of the pulsar
needs to be ∼1000 km s−1 . Although 1000 km s−1 is somewhat
large as a pulsar kick velocity, it is still within the acceptable
range (e.g., Ng & Romani 2007), and the pulsar in another Orich SNR Puppis A may also have such a high velocity (Hui &Becker 2006).
4.4.2. Ejecta Distribution and Explosion Asymmetry

The optical and X-ray studies of the SN ejecta in G292.0+1.8
have shown that their spatial distribution and physical properties
are not symmetric, including the northwest–southeast concentration of the O and Ne ejecta, and the distinctive differenceof the X-ray plasma temperature between the northwestern andsoutheastern areas (Park et al. 2002, 2007; Ghavamian et al.
2005). Also the ejecta in the northern and southern boundaries
move faster than those in the eastern and western areas (Winkler et al. 2009). It is worth noting that the pulsar jet axis is also
along the northeast–southwest direction (Park et al. 2007).
Our AKARI result on the Ne-line emitting ejecta material
identified by their high L15/L24 ratio is consistent with thespatial distribution seen in [O iii] 5007: most of the Ne+ /Ne++
ions are distributed in the southeastern area called spur andstreamers; several isolated ones coincide with the optical knots,
and the other group of Ne-rich knots is distributed in thenorthwestern area and coincides with [O iii] optical emission
knots (fast-moving knots, or FMKs, as reported by Ghavamian
et al. 2005 and Winkler & Long 2006). The northwest ejecta isless obvious in optical but it is easily recognized in the X-ray
image of ionized O and Ne elements (Park et al. 2002). This
type of bipolar ejecta distribution is also found in other young
core-collapse SNRs, e.g., Cas A and G11.2−0.3 (Smith et al.
2009; Koo et al. 2007; Moon et al. 2009). Cas A shows a similar
Ne ejecta distribution, which is almost perpendicular to the wellknown narrow northeast–southwest jet, but is roughly aligned
to the bipolar ionic ejecta, which sometimes is suggested be
the major direction of explosion (Hwang et al. 2004; Wheeler
et al. 2008; Smith et al. 2009). In G11.2−0.3, which is theremnant of the historical SN AD 386, the iron ejecta is found to
be distributed mainly along the northwest–southeast direction(Koo et al. 2007; Moon et al. 2009).

Vol. 706

The symmetry axis in the spatial and kinematical distributionof ejecta in G292.0+1.8 is, therefore, either along northwest–
southeast or north–south, which is perpendicular to the plane of
the ER. We consider that either the explosion was asymmetricand/or the CS wind was denser in the equatorial plane so that theejecta expanding in this plane was slowed down more comparedto those expanding to the other directions.
An interesting feature is the narrow tail in the wide L18Wband image that extends from the end of the streamers to outside
the remnant. A possible explanation is that the narrow tailis a part of the O-rich streamer. Note that the narrow tail isconnected to the streamers by a southern patch of emission near
the boundary of L15-band image. The distance from the center
to the end of the narrow tail is 7 (13 pc), which is ∼2 (3 pc)
farther than the outermost O-rich clump in this area (Winkler
et al. 2009). It, however, does not has an optical counterpart
in the [O iii] image (Winkler et al. 2009). If we assume aconstant velocity and adopt an age of 3000 years (Ghavamian
et al. 2005; Winkler et al. 2009), the transverse velocity of thenarrow tail is 4000 km s−1 . This gives a possibility that its radial
velocity is also very large, beyond the velocity range of previous
optical narrow-band imaging observations (e.g., ∼2000 km s−1
in Winkler & Long 2006). An alternative explanation is that thenarrow tail is a reradiated IR light echo similar to the echoes
identified in Cas A (Krause et al. 2005; Dwek & Arendt 2008).
The lack of counterpart in the optical [O iii] and X-ray images,
together with the faint feature in the FIR images (Figure 6), could
be consistent with the continuum origin of the IR emission. In
the case of a light echo, the location of the echo can be derived
from the geometrical equation of an ellipse whose two focii are
the SN and the observer (e.g., Couderc 1939; Dwek & Arendt2008). Applying a projected radius of 13 pc and a time delayof 3000 years, we obtain a location of the IR echo at 450 pc
behind G292.0+1.8 and its angular offset of ∼2◦ from the lineof sight. It appears that this alignment is too tight at a first glance,
but it can be a selection effect caused by our limited imaging
area. For example, Krause et al. (2005) discovered IR echoes
along the scan direction in Cas A for the first time, but Dwek &Arendt (2008) reported that the echoes were distributed in many
positions in a large (∼2◦ ) area. More observations are definitely
necessary to inspect the nature of the narrow tail.
On the other hand, the sharp boundary of the SN blast wave
was detected only toward the north and southwest (Figure 2).
The absence of the SN blast wave in the other directions implies
that either the shock is trapped within the shell because thatpart of the shell has a larger column density or the SN blast
wave has propagated far beyond because the ambient density
is lower toward that direction. There is no indication in the
AKARI images that the OES is denser, where the SN blast wave
is missing. Instead those parts are fainter in the FIR images
which indicate a lower column density. In this regard, the faintMIR emission that extends far beyond the bright shell to thesoutheast (Figure 5) is interesting because if it is part of the SNR,
it indicates that the SN blast wave has propagated much further
out toward this direction probably due to the lower ambientdensity. A deep radio observation could reveal faint featuresassociated with this IR structure.
5. CONCLUSIONS
We have presented the NIR to FIR imaging and MIR
spectroscopic observations of the O-rich SNR G292.0+1.8 using
the IRC and FIS instruments aboard AKARI satellite. The almost
continuous multi-band imaging capability of AKARI coveringNo. 1, 2009AKARI INFRARED OBSERVATIONS OF THE SNR G292.0+1.8

wide IR wavelengths together with its wide FoV enabled us toclearly see the distinct IR emission from the entire SNR and to
derive its IR characteristics. We derive the physical parametersof IR-emitting dust grains in the swept-up CSM and compare theresult with the model calculations of dust destruction by a SN
shock. The overall shape of the observed SED can be explainedby a simple model using characteristic SNR parameters witha bit lower initial dust-to-gas ratio. At 11 μm, the model flux
is significantly smaller, which may indicate the importance ofstochastic heating. We have not detected any significant amountof freshly formed dust associated with the SN ejecta.
The AKARI results in this paper give new insight into the
explosion dynamics of G292.0+1.8. We have discovered an almost symmetric IR shell probably produced by the circumstellar wind from the progenitor star in the RSG phase. Its centeris significantly offset from the previously suggested explosion
centers. We consider that either the OES represents the circumstellar shell pressure-confined by an external medium or
the SN exploded close to the center of the OES. In the latter
case, the pulsar in G292.0+1.8 may be traveling at a speed of∼1000 km s−1 . At the same time, the ejecta distribution is unveiled by their high 15/24 μm ratio. The ejecta are mainly
distributed along the northwest–southeast direction. This symmetric pre-SN structure and asymmetric ejecta distribution appear to be rather common in the remnants of SN IIL/b, whichsuffer strong mass loss like Cas A (Hines et al. 2004; Smith et al.
2009). There is also a narrow tail outside the SNR shell, whichmight be similar to the feature observed in Cas A. A detailed
study is necessary to understand the nature of this feature.
We suggest that multi-band IR imaging observations are
powerful tools to explore both the ejecta and CSM emission inyoung core-collapse SNRs. Especially, the 15 and 24 μm images
are useful to reveal the detailed structure of IR features, whichleads to better understanding of environments of the progenitor
and the SN explosion.
This work is based on observations with AKARI, a JAXA
project with the participation of ESA. We thank all the membersof the AKARI project. We thank an anonymous referee forcareful reading and constructive comments. We also thankB. Gaensler for providing the ATCA 20 cm image and S.
Park for providing the Chandra X-ray image. This work was
supported by the Korea Research Foundation Grant fundedby the Korean Government (KRF-2008-357-C00052) and the
Korea Science and Engineering Foundation (R01-2007-00020336-0). This work was also supported in part by the Grant-inAid for Scientific Research for the Japan Society of Promotion of
Science (18204014, 20340038). T.N. has been supported in partby the World Premier International Research Center Initiative(WPI Initiative), MEXT, Japan, and by the Grant-in-Aid for
Scientific Research of the Japan Society for the Promotion of
Science (19740094).
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