The Astrophysical Journal, 749:173 (12pp), 2012 April 20

doi:10.1088/0004-637X/749/2/173


C 2012. The American Astronomical Society. All rights reserved. Printed in the U.S.A.

A SEARCH FOR INFRARED EMISSION FROM CORE-COLLAPSE SUPERNOVAE
AT THE TRANSITIONAL PHASE
Masaomi Tanaka1,2 , Takaya Nozawa2 , Itsuki Sakon3 , Takashi Onaka3 , Ko Arimatsu3 , Ryo Ohsawa3 ,Keiichi Maeda2 , Takehiko Wada4 , Hideo Matsuhara4 , and Hidehiro Kaneda5
1 National Astronomical Observatory of Japan, Mitaka, Tokyo, Japan; masaomi.tanaka@nao.ac.jp
2 Institute for the Physics and Mathematics of the Universe, University of Tokyo, Kashiwa, Japan; takaya.nozawa@ipmu.jp
3 Department of Astronomy, Graduate School of Science, University of Tokyo, Bunkyo-ku, Tokyo, Japan
4 Institute of Space and Astronautical Science, Japan Aerospace Exploration Agency, Sagamihara, Kanagawa, Japan
5 Graduate School of Science, Nagoya University, Chikusa-ku, Nagoya, Japan
Received 2011 January 11; accepted 2012 February 17; published 2012 April 5

ABSTRACT
Most of the observational studies of supernova (SN) explosions are limited to early phases (<a few years after the
explosion) of extragalactic SNe and observations of SN remnants (>100 yr) in our Galaxy or very nearby galaxies.
SNe at the epoch between these two, which we call the “transitional” phase, have not been explored in detail
except for several extragalactic SNe including SN 1987A in the Large Magellanic Cloud. We present theoreticalpredictions for the infrared (IR) dust emissions by several mechanisms; emission from dust formed in the SN ejecta,
light echo by circumstellar (CS) and interstellar (IS) dust, and emission from shocked CS dust. We search for IR
emission from six core-collapse SNe at the transitional phase in the nearby galaxies NGC 1313, NGC 6946, andM101 by using the data taken with the AKARI satellite and Spitzer. Among six targets, we detect the emission
from SN 1978K in NGC 1313. SN 1978K is associated with 1.3 × 10−3 M of silicate dust. We show that, among
several mechanisms, the shocked CS dust is the most probable emission source to explain the IR emission observedfor SN 1978K. IR emission from the other five objects is not detected. Our current observations are sensitive to IR
luminosity of >1038 erg s−1 , and the non-detection of SN 1962M excludes the existence of the shocked CS dust fora high gas mass-loss rate of ∼10−4 M yr−1 . Observations of SNe at the transitional phase with future IR satellites
will fill the gap of IR observations of SNe with the age of 10–100 yr, and give a new opportunity to study the CS
and IS environments of the progenitor, and possibly dust formation in SNe.
Key words: dust, extinction – infrared: stars – supernovae: general – supernovae: individual (SN 1978K)
Online-only material: color figures

This transitional phase is important to understand a longterm evolution of SNe, from SN to SNR. A recent study by
Larsson et al. (2011) shows that SN 1987A, in fact, experiencesa transition from SN to SNR in this phase. However, sincemost observations of SNe at the transitional phase except for
SN 1987A have been performed at X-ray or radio wavelengths,
the entire properties of SNe at the transitional phase are not fully
understood.
In this paper, we present the results of our search for infrared(IR) emission from core-collapse SNe at the transitional phase.
IR emission from SNe is expected to arise from dust associatedwith SNe. There are three possible populations of dust grains,according to their locations: (1) dust formed in the ejecta of SNe(SN dust), (2) dust formed by the pre-SN stellar wind, which
is now located in the circumstellar (CS) region (CS dust), and(3) interstellar dust (IS dust). CS and IS dust can contributeto the IR emission by absorbing the radiation from SNe andreemitting it (light echo). CS dust is also heated and can also
emit IR emission when it is swept up by the SN shock. In theseways, dust grains in and around SNe convert the large kinetic
energy and luminosity of the SNe into the IR luminosity. Thus,
IR observations of SNe can be used as a probe of the progenitor
environment.
In Section 2, we describe the target selection, the data usedin the analysis, and results of photometry. We use the data takenwith two IR satellites, AKARI satellite (Murakami et al. 2007)
and Spitzer (Werner et al. 2004), which provide the deepest
images in near-IR (NIR) and mid-IR (MIR) wavelengths. Dust

1. INTRODUCTION
Core-collapse supernovae (SNe) are the explosions of massive stars at the end of their lives. SN ejecta expand into theinterstellar medium (ISM) or circumstellar medium (CSM) witha huge kinetic energy (∼1051 erg). First, SN ejecta experience a
free expansion. As the ejecta sweep up the surrounding material,
the SN begins to decelerate. When the swept-up mass becomescomparable to the mass of the SN ejecta at t ∼ a few hundreds
years (hereafter t denotes the time after the explosion), the SN
ejecta expand self-similarly (so-called Sedov phase). Finally,
the SN merges into the ISM at t  10,000 yr.
Observational studies of SNe have been performed mainly inthe two distinct ways. One is the observations of extragalactic
SNe at their early phases (t < 10 yr), and the other isobservations of supernova remnants (SNRs, t  100 yr) in
our Galaxy or very nearby galaxies. In contrast, SNe at theepoch between these two (t = 10–100 yr), which we call the“transitional” phase, have not been explored in detail. At such
epochs, only a few objects have been studied. The best case is
SN 1987A in the Large Magellanic Cloud (McCray 2007 andreferences therein). Other examples include several X-ray SNe(e.g., Immler & Kuntz 2005; Immler et al. 2005; Soria & Perna2008), and long-lasting SNe, such as SNe 1978K (Ryder et al.
1993; Schlegel et al. 1999, 2000, 2004; Lenz & Schlegel 2007;
Smith et al. 2007) and 1988Z (Williams et al. 2002; Schlegel &Petre 2006). See also a recent paper on SN 1980K by Sugerman
et al. (2012).
1

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Table 1
List of Targets

Host Galaxy

Distance

NGC 1313

4.13 ± 0.11

NGC 6946

5.6 ± 0.3M101

6.7 ± 0.3

SN

SN Type

SN 1962M
SN 1978K
SN 1917A
SN 1968D
SN 1909Aa
SN 1951H

II
IIn
II
II
II peculiar
II

(Mpc)

Position

Position Ref.

α (J2000.0)

δ (J2000.0)

03h 18m 12.s 2
03h 17m 39.s 0
20h 34m 46.s 90
20h 34m 58.s 41
14h 02m 03.s 1
14h 03m 55.s 3−66◦ 31 38−66◦ 33 04
+60◦ 07 29. 08+60◦ 09 34. 48+54◦ 27 58+54◦ 21 41

1
2
3
4
3
3

References. (1) Artyukhina et al. 1996; (2) Dopita & Ryder 1990; (3) Barbon et al. 1999; (4) van Dyk et al. 1994.
a SN 1909A is out of field of view in the IRC images, and only in the Spitzer/MIPS image.

Each channel has about a 10 × 10 field of view. The NIR and
MIR-S share the same field of view while the MIR-L observes a
sky about 25 away. Thus, the data we use in this paper consist of
two observational sequences. The pixel scales of the detectors
are 1. 46 × 1. 46, 2. 34 × 2. 34, and 2. 51 × 2. 39 for the NIR,
MIR-S, and MIR-L, respectively.
We use the archived imaging data taken in the AKARI
mission program “ISM in our Galaxy and Nearby galaxies”(ISMGN; Kaneda et al. 2009). The data are taken with thetwo-filter mode (Astronomical Observation Template (AOT)
IRC02; see Onaka et al. 2007). We use six band images intotal, i.e., N3 (reference wavelength 3.2 μm), N4 (4.1 μm), S7(7.0 μm), S11 (11.0 μm), L15 (15.0 μm), and L24 (24.0 μm).
The wavelength coverage with the response larger than 1/e ofthe peak is 2.7–3.8 μm, 3.6–5.3 μm, 5.9–8.4 μm, 8.5–13.1 μm,
12.6–19.4 μm, and 20.3–26.5 μm for N3, N4, S7, S11, L15, andL24, respectively. For more details of the IRC instrument, see
Onaka et al. (2007). The log of observations is summarized inTable 2.
The data were reduced by using the IRC imaging pipeline
software version 20091022 (see Lorente et al. 2007 for details).
The pipeline performs the basic reduction, e.g., correction ofbad pixels, subtraction of dark current, rejection of cosmic rays,
linearity correction, flat fielding, and co-adding of the frames.
To determine the accurate position, we match the sources in
the N3, N4, and S7 band images with those in Two Micron All
Sky Survey catalog (Skrutskie et al. 2006). The uncertainty inthe position is <2 for the N3 and N4 band images, and <3
for the S7 band images. Then, the S11, L15, and L24 images
are matched with S7 images. In these images, the uncertainty inposition is estimated as ∼3 , which is larger than that in the N3,
N4, and S7 images because the matching is based on a small
number of common sources in the field of view.
To obtain the flux densities of the sources, we perform aperture photometry with the aperture of 10 pixels and 7.5 pixels for
the NIR and MIR-S/L bands, respectively, As the background,
we use the annulus of 5 pixel width just outside of the aperture
radius. When the SN is located at the edge of the image, the aperture correction is applied. The flux is calibrated by following
Lorente et al. (2007). Color correction is not performed to thederived photometry because the intrinsic spectral shape of thesource is not known. Instead, we apply correction to the model
when it is compared with the observations (see Section 3).
When the SN is not detected, we derive an upper limit of theflux by putting an artificial source and measuring the flux of thesource. For this purpose, we need a point-spread function (PSF)
of each image. For the N3 and N4 images, PSF is constructed by
the point sources in the same image. Since there are not manymodels used in this paper are shown in Section 3. Among six
targets, we detect IR emission from SN 1978K. The originof this emission is discussed in Section 4. In Section 5,
theoretical predictions for the IR emission from normal SNeat the transitional phase are summarized, and implications of
non-detection of five SNe are discussed. Prospects for future
observations are discussed in Section 6. Finally, we summarize
the conclusions in Section 7.
2. TARGETS AND OBSERVATIONS
2.1. Targets
To search for IR emission from SNe at the transitional phase,
we target seven core-collapse SNe, namely, SNe 1909A, 1917A,
1951H, 1962M, 1968D, and 1978K, discovered in three very
nearby galaxies NGC 1313, NGC 6946, and M101. These targets
are selected based on their proximity to the host galaxy and theavailability of NIR and MIR imaging data taken with the Infrared
Camera (IRC; Onaka et al. 2007) onboard the AKARI satellite(Murakami et al. 2007). General information on our targets issummarized in Table 1.
Note that SN 1948B (Type II) in NGC 6946 is excluded
because of the uncertainty in its position. SNe 1939C and 1969P
in NGC 6946 are also excluded since SN 1939C is classified asmerely Type I (either core-collapse SNe or Type Ia SNe) andthe type of SN 1969P is unknown. SN 1980K is also locatedin NGC 6946. But since it is found that AKARI data are notas deep as the data presented by Sugerman et al. (2012), we
do not include SN 1980K. SN 1970G (Type IIL) in M101 is
excluded because the SN position is heavily contaminated by
the H ii region NGC 5455 (Fesen 1993).
For the distance to NGC 1313, we adopt 4.13 ± 0.11 Mpc,
which is estimated from the tip of the red giant branch (Méndez
et al. 2002). We use this value since it was also used by Lenz
& Schlegel (2007) and Smith et al. (2007) for the study of
SN 1978K. For NGC 6946, there are several distance measurements. We adopt the average value of 5.6 ± 0.3 Mpc from the
H i Tully–Fisher relation (Pierce 1994), the CO Tully–Fisher
relation (Schoniger & Sofue 1994), and the expanding photosphere method of Type II SN 1980K (Schmidt et al. 1994)
and SN 2004et (Sahu et al. 2006). The distance to M101 is assumed to be 6.7 ± 0.3 Mpc, which is estimated by the Cepheidcalibration (Freedman et al. 2001).
2.2. AKARI IRC Data
The AKARI/IRC equips three channels, i.e., NIR (1.8–5.5 μm), MIR-S (4.6–13.4 μm), and MIR-L (12.6–26.5 μm).
2

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Table 2Log of AKARI/IRC Observations

GalaxyNGC 1313
NGC 6946
M101

Date

Observation ID

Exposure Time (s)NIR, MIR-S

MIR-L

NIR, MIR-S

MIR-L

NIR

MIR-S, L2006 Dec 32006 Dec 18
2007 Jun 14

2006 Dec 42007 Jun 16
2007 Jun 17

1400416-001
1400620-001
1402211-001

1400417-001
1402217-001
1402212-001

178
178
178

196
196
196

Table 3
Summary of Spitzer/MIPS Data
GalaxyNGC 1313
NGC 6946
M101

Date

Exposure Time(s)

Observing ModeAOR key

Program ID

2007 Sep 29
2007 Jul 102008 Jan 5

42
493
42

Medium scan
Photometry
Medium scan

22618112
18271232
21380144

40204: Kennicut et al.
30494: Sugerman et al.
40352: Rieke et al.

3. DUST MODELS

sources in the S7, S11, L15, and L24 images, we use an average
PSF constructed from sources in different fields (Arimatsu et al.
2011).

In the following sections, we compare the observations withvarious model spectra. We consider two kinds of carbon dust
grains, i.e., amorphous carbon and graphite, and two kinds ofsilicate dust grains, i.e., “astronomical silicate” and forsterite(Mg2 SiO4 ). We assume that grains are spherical.
Thermal emission flux from the dust toward an SN can be
expressed as follows:
4π 3 −11, (1)ρaFλ = Mdust
4π a 2 Qabs
π
B
(T
)
λ
dust
λ
3
4π d 22.3. Spitzer MIPS Data
In addition to the AKARI/IRC data, we use archive data takenwith the Multiband Imaging Photometer for Spitzer (MIPS;
Rieke et al. 2004) onboard Spitzer (Werner et al. 2004).
All the six SNe are within the field of view of someobservations with the MIPS 24 μm band. NGC 1313 wasobserved with the scan mode in the Spitzer Local Volume
Legacy program (Dale et al. 2009). M101 was also observedwith the scan mode by Gordon et al. (2008). SNe 1917A and1968D in NGC 6946 are within the field of SN 2002hh (Meikle
et al. 2006). A summary of the data, including the Astronomical
Observation Request (AOR) keys, program IDs, and PI names,
is shown in Table 3.
We use the post-basic calibrated data (PBCD) products, which
are produced by the Spitzer pipeline. The pixel scale of the
final image is 2. 45 × 2. 45. The reference wavelength of the
MIPS 24 μm band is 23.68 μm and the wavelength coverage is
20.5–28.5 μm (the range where the response is >10% of thepeak). The positional uncertainty is 1. 4.
For photometry, we perform aperture photometry. When the
SN is not detected, we derive an upper limit of the flux by putting
an artificial source around the SN position. We constructed a PSF
from point sources in the scan mode images of NGC 1313 andM101 and use it to derive the upper limits.
The Spitzer/MIPS data are found to be slightly deeper thanthe AKARI/IRC data at 24 μm (Section 2.4). Thus, we use MIPS
data for the flux of 24 μm.

where Mdust is the total dust mass, ρ and a are, respectively, thedensity and the radius of the dust grain, Qabs
λ is the absorption
efficiency factor, Tdust is the temperature of the dust, and d isthe distance to the SN. We assume a single temperature forsimplicity. In the equation above, all of the IR emission fromdust grains is assumed to be detected without being absorbed.
We consider a single size of dust a = 0.1 μm. In reality, there
is a distribution in the grain size, and if dust grains are swept up
by the shock waves of the SN, their size distribution can changewith time (Nozawa et al. 2007). However, Qabs
λ /a is independentof the grain size if λ  2π a, and thus, the assumption of thesingle-size dust does not affect the estimate of the dust mass.
For the grain densities, we adopt ρ = 2.26, 2.26, 3.8,
and 3.2 g cm−3 for amorphous carbon, graphite, astronomicalsilicate, and forsterite, respectively. The absorption efficiency
factor is taken from Edoh (1983) for amorphous carbon, Draine(2003b) for graphite and astronomical silicate, and Semenov
et al. (2003) for forsterite.
Since the observed flux in Table 4 is not color-corrected dueto the unknown intrinsic spectrum of the sources, we apply
the color corrections to the model. Since the flux density of the
AKARI/IRC data refers to the spectrum Fλ ∝ 1/λ, the following
correction is applied to the model spectrum Fλmodel (λ):
 modelF(λ)R(λ)dλ.
(2)
FλIRC (λi ) =  λ(λi /λ)R(λ)dλ2.4. Results of Photometry
We detect IR emission associated with SN 1978K in all the
bands (Table 4). Figure 1 shows the AKARI/IRC images aroundSN 1978K (180 × 180 section). The source has an extremely
red color with a flux ratio of Fν (S7)/Fν (S11) = 0.1, which isout of a typical range in the spiral arm and the inter-arm of
the galaxy (∼0.8–1.4; Sakon et al. 2007). The origin of thisemission is discussed in Section 4.
For the other five objects, we do not detect a significant signal
from SNe. Upper limits of the flux are shown in Table 4.

Here, λi is the reference wavelength of the band and R(λ) is
the response function (electron energy−1 ). We use the response
function measured in the laboratory6 (see Onaka et al. 2007;63

http://www.ir.isas.jaxa.jp/AKARI/Observation/

The Astrophysical Journal, 749:173 (12pp), 2012 April 20

Tanaka et al.

Figure 1. AKARI/IRC images around SN 1978K in NGC 1313 (180 × 180 section, which corresponds to about 3.6 kpc × 3.6 kpc). North is up, and east is tothe left. Note that the PSF of the N3 and N4 bands are slightly elongated in images on orbit due to a misalignment of the telescope and color aberrations in the NIR
channel (Kaneda et al. 2007).

Table 4
Summary of Photometry
SN

Agea
(yr)

IRC/N3

IRC/N4IRC/S7

IRC/S11

IRC/L15

IRC/L24

MIPS/24μm

SN 1962M

44.0
44.9
28.4
29.2
89.5
90.0
90.0
38.8
39.3
39.4
99.0
55.9
56.3

<0.2
...
0.074 (0.014)
...
<0.1
...
...
<0.3
...
...
...
<0.07
...

<0.2
...
0.086 (0.010)
...
<0.2
...
...
<0.3
...
...
...
<0.1
...

<0.2
...
0.287 (0.039)
...
<8
...
...
<4
...
...
...
<0.4
...

<0.2
...
2.85 (0.24)
...
<6
...
...
<5
...
...
...
<0.6
...

<0.9
...
2.53 (0.13)
...
...
<10
...
...
<6
...
...
<1
...

<2
...
2.73 (0.20)
...
...
<10
...
...
<8
...
...
<3
...

...
<1
...
3.06 (0.40)
...
...
<10
...
...
<5
<0.7
...
<1

SN 1978K
SN 1917A

SN 1968D

SN 1909A
SN 1951H

Flux (mJy)b

Notes.
a Age is given as year after the discovery, which is a sound approximation for the age after the explosion.
b Values in parenthesis for SN 1978K represent 1σ error including the measurement error and the calibration error (5 %). In the case of non-detection, 3σ upper limit
is given.

Lorente et al. 2007). The calibration of the Spitzer/MIPS datauses the blackbody spectrum with T0 = 10,000 K as thereference. Thus, we apply the following correction:
 modelF(λ)R(λ)dλ
MIPS
. (3)
Fλ (λi ) =  5 λ
 λi
exp(hc/λi kT0 ) − 1
R(λ)dλλexp(hc/λkT0 ) − 1

4. IR EMISSION FROM SN 1978K
4.1. Spectral Fit
Among our seven targets, we detect IR emission only fromSN 1978K. SN 1978K is an extraordinarily strong Type IIn SN,
which is visible both in radio and X-ray even at ∼30 yr after
the explosion (Ryder et al. 1993; Petre et al. 1994; Chugai et al.
1995; Schlegel et al. 1996, 1999, 2000, 2004; Montes et al. 1997;
Chu et al. 1999; Lenz & Schlegel 2007; Smith et al. 2007). Suchlong-lasting radio and X-ray emissions are clear evidence forthe presence of the dense CSM, with which the SN ejecta has
been interacting (Chevalier 1982; Chevalier & Fransson 1994;
Chugai et al. 1995).

We use the response function provided by the Spitzer Science
Center.7
7

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Tanaka et al.

101

102

SN 1978K

SN 1978K
101

-0.60

ν

0

Flux (mJy)

Flux (mJy)

10
0

-0.67

ν

10

-1

10

Amorphous carbon
Tdust = 180 K-3
Mdust = 6.8 x 10 M

10-1

10-2

10-3

Cas A
(at 4.13Mpc)
1

10

10-2

100
1000 10000 100000
Wavelength (μm)

Astronomical silicate
Tdust = 230 K
Mdust = 1.3 x 10-3 M
1

10

100

Wavelength (μm)
Figure 3. IR SED of SN 1978K (open black circle) compared with the model
spectra. The red line shows the spectrum of the amorphous carbon with Tdust =180 K and Mdust = 6.8 × 10−3 M while the blue line shows the spectrumof astronomical silicate with Tdust = 230 K and Mdust = 1.3 × 10−3 M . The
crosses show the color-corrected model flux for each band. For both models,
the synchrotron component (Fν ∝ ν −0.60 ) is added. In the case of astronomicalsilicate, the color corrections are FνIRC,MIPS (λi )/Fνmodel (λi ) = 1.7, 0.91, 1.1,
and 1.1 for the S7, S11, L15, and MIPS 24 μm band, respectively. In the caseof amorphous carbon, these are 1.1, 0.93, 0.95, and 1.0, respectively.
(A color version of this figure is available in the online journal.)

Figure 2. IR–radio SED of SN 1978K (filled circles) compared with that of Cas
A (open squares) at the same distance (4.13 Mpc). The red dashed line shows
the power-law fit to the IR and radio data of SN 1978K while the blue dashedline shows the fit only to the radio data. The radio data of SN 1978K are takenfrom Smith et al. (2007). The data of Cas A are taken from Baars et al. (1977),
Mezger et al. (1986), Liszt & Lucas (1999), Hines et al. (2004), and Barlow
et al. (2010).
(A color version of this figure is available in the online journal.)

Figure 2 shows the spectral energy distribution (SED) of
SN 1978K from IR to radio wavelengths (filled circles). The
SED of Cassiopeia A (Cas A) at the distance of NGC 1313 isalso shown for comparison (open squares). These two objects
show similar overall SEDs from IR to radio wavelengths.
Since SN 1978K is powerful in radio wavelengths (Smithet al. 2007), we must take into account the influence of thesynchrotron emission on the IR bands (see, e.g., Mezger et al.
1986; Rho et al. 2003; Dwek 2004; Hines et al. 2004; Barlow
et al. 2010, for the case of Cas A). The radio emission of
SN 1978K is well fitted by a power-law spectrum Fν ∝ ν α(Montes et al. 1997; Smith et al. 2007). When we fit the flux offour radio bands of SN 1978K (Smith et al. 2007), we derive
α = −0.67 ± 0.05 (blue dashed line in Figure 2). The flux
densities at the N3 and N4 bands are almost in line with thepower-law spectrum. If we include the flux of these two bands
for the fitting, we obtain α = −0.60 ± 0.02 (red dashed line).
Since these two values are consistent within the error, we assume
that the emissions at the N3 and N4 bands are of synchrotron
origin and use α = −0.60 in the following discussions.
The flux density of the S7, S11, L15, and L24 (MIPS 24 μm)
bands is clearly above the synchrotron component. Therefore,
we conclude that this excess is caused by the dust emission
associated with SN 1978K.
We fit the observed flux with the model spectra of the four
dust species with the dust temperature and mass as parameters.
The quality of the fit is judged by χ 2 for the four observational
points, i.e., the IRC S7, S11, L15, and MIPS 24 μm bands.
For each dust species, we choose the best model that gives thesmallest χ 2 . Since the dependence on the temperature and mass
is simple, the fit is unique, i.e., there is no other solution with adifferent set of temperature and mass.
Among the four dust species, the smallest χ 2 is obtained
for astronomical silicate with Tdust = 230 K and Mdust =1.3 × 10−3 M (blue line in Figure 3). The total IR flux by thedust emission is FIR = 7.0 × 10−13 erg cm−2 s−1 and the total
IR luminosity is LIR = 1.5 × 1039 erg s−1 . Silicate dust grains

are favored over carbon grains because the observation shows
excess at the S11 band, which can be attributed to the Si–O
stretching mode of silicate dust at 10 μm (Draine & Lee 1984;
Draine 2003a). As shown in the red line in Figure 3, amorphous
carbon dust grains are not able to explain the excess. Forsterite
fits the observations quite well, but gives a band feature at
10 μm slightly narrower than astronomical silicate. Hereafter,
we simply call the dust species found in SN 1978K “silicate
dust.”
4.2. Origin of the IR Emission
SN 1978K is associated with 1.3 × 10−3 M of silicate dust.
The derived temperature is relatively high (230 K). In thissubsection, we discuss possible origins of the IR emission. Since
SN 1978K is a CSM-rich SN, we first consider CS dust. CS dust
can radiate IR emission by light echo (Section 4.2.1). Then,the echoing CS dust is located at the distance larger than D =ct/2 = 1.4 × 1019 (t/30 yr) cm, where the SN forward shock
cannot reach. On the other hand, the CS dust swept up by the SN
forward shock can be heated via the collisions with energetic
electrons and emits thermal radiation (Section 4.2.2). Then, thetypical location is at ∼5 × 1017 (t/30 yr)(vSN /5000 km s−1 ) cm.
Another possibility is the SN dust (Section 4.2.3). We show
that, among these possibilities, the shocked CS dust is the mostprobable origin for the IR emission from SN 1978K.
4.2.1. Light Echo

The first possibility for the IR emission is echo emission of
CS dust. The dust can absorb the ultraviolet or optical photons
of an SN, and radiate IR emission (Bode & Evans 1980; Wright1980; Dwek 1983, 1985; Graham & Meikle 1986). In fact, thepresence of CS dust echo has been suggested by early phaseobservations of SNe (Meikle et al. 2006, 2007; Andrews et al.
2011) and SN 1980K at the transitional phase (Sugerman et al.
2012).
5

The Astrophysical Journal, 749:173 (12pp), 2012 April 20

Luminosity (erg s-1)

CS echo
1041
1040

1042

SN 1978K IR
SN 1978K Opt B
SN 1978K Opt R
Lin
LIR

IS echoLuminosity (erg s-1)

1042

Tanaka et al.

1039
10

38
5

10

twind = 106 yr

0

5

10
15
Time (yr)

20

25

1040
1039

n = 1 cm-3
10

38

10

37

n = 0.1 cm

twind = 10 yr

37

1041

SN 1978K IR
SN 1978K Opt B
SN 1978K Opt R
Lin
LIR

30

0

5

10
15
Time (yr)

20

25

-3

30

Figure 4. Left: computed light curves of the IR echo by the CS dust (lines), compared with observations of SN 1978K (points). The blue line shows the input
SN luminosity, while the black line shows the IR luminosity by the light echo. The optical data points are shown by assuming zero bolometric correction using
the data by Ryder et al. (1993) and Schlegel et al. (1999). The IR echo of the SN outburst luminosity L = L0 exp(t/tSN ). The solid line shows the model with
L0 = 6 × 1042 erg s−1 , tSN = 25 days, Ṁ = 10−4 M yr−1 , and twind = 105 yr. Dashed lines show the extreme cases with twind = 106 yr. Even with such an extreme
assumption, the expected luminosity is lower than the observed IR luminosity at t = 30 yr. Right: the same with the left panel but for the IS dust echo. Black solid anddashed line shows the IR luminosity for n0 = 0.1 and 1 cm−3 , respectively.
(A color version of this figure is available in the online journal.)

with the radius of R2 = 100 pc (t = 2R2 /c ∼ 700 yr). The
black solid and dashed lines in the right panel of Figure 4 show
the luminosity evolution of the IS dust echo for the ISM gas
density of n = 0.1 and 1 cm−3 , respectively.
With the normal SN luminosity, the expected IR luminosity at
t = 30 yr is not high enough to explain the observed luminosityof SN 1978K. To explain the observation, the SN had to bevery luminous, L0 ∼ 6 × 1043 erg s−1 , or the absolute opticalmagnitude of about −20 mag. It would imply that the brightnessof SN 1978K was of the 8th magnitude. This is in contrast
to the suggestion by Ryder et al. (1993) that SN 1978K was
subluminous. Although a very large luminosity of SN 1978K is
not completely excluded given the sparse observations in 1978,
we conclude that IS dust echo is also unlikely.

We calculate the luminosity of the light echo by following
Dwek (1983). The gas mass-loss rate of the progenitor of
SN 1978K is estimated to be about Ṁ = 10−4 M yr−1
from optical, X-ray, and radio observations (Ryder et al. 1993;
Chugai et al. 1995; Smith et al. 2007). The duration of massloss is assumed to be twind = 105 yr. This is a rather extreme
assumption since the progenitor loses mass as large as 10 M
prior to the explosion. We assume that the wind velocity is
vwind = 10 km s−1 , which gives outer radius of CSM about
R2 = 3×1018 cm. The mass-loss rate is assumed to be constant,and thus, the CSM density has a power-law radial profile ∝ r −2 .
For the SN luminosity, we use an exponential form L =L0 exp(t/tSN ). We assume a normal peak luminosity L0 =6 × 1042 erg s−1 and tSN = 25 days (Sahu et al. 2006; Kotak
et al. 2009). The other parameters are set as follows: the gas-todust mass ratio f = 200, the grain radius a = 0.1 μm, the mean
absorption efficiency Q̄ = 1, and the evaporation temperatureof dust T = 1500 K. The SN outburst luminosity will make a
dust cavity with a radius of R1 = 7 × 1016 cm.
The left panel of Figure 4 shows the computed IR luminosityof the CS dust echo (solid black line). Optical data of SN 1978K
are also plotted for comparison under a simple assumption ofzero bolometric correction (blue and red points). The IR echo
luminosity at t = 10–100 yr depends largely on the outerradius (R2 ) of the CSM since the luminosity from the CS echo
declines rapidly at t = 2R2 /c ∼ 7 yr (Dwek 1983). To illustrate
this dependence, we also show the model with a longer mass-loss
duration (dashed line for twind = 106 yr). With twind  5×105 yr,
with which the progenitor loses more than 50 M , the IR echodoes not drop dramatically at t < 30 yr. But, even with thisextreme case, the IR luminosity is only an order of 1037 erg s−1
at t = 30 yr, which is much less than the observed IR luminosity.
Thus, we conclude that the CS dust echo of SN outburst isunlikely to be the origin of the IR emission.
We also calculate the IS dust echo by the same method as thatfor the CS dust echo. In fact, IR echo by the IS dust has been
detected for several extragalactic SNe (Meikle et al. 2007, 2011;
Kotak et al. 2009) and Cas A (Krause et al. 2005; Dwek & Arendt2008). The difference from the CS dust is that the density is set
to be uniform and that the much larger outer radius is allowed.
Here, the distribution of IS dust is approximated to be a sphere4.2.2. Shocked CS Dust

The second source for the IR emission is the CS dust heatedby the forward shock of the SN. The mid-IR emission from thering in SN 1987A is thought to arise from this mechanism (e.g.,
Dwek et al. 2008). It is interesting to note that the emission
properties of SN 1987A ring and SN 1978K look similar. The
dust temperature in the SN 1987A ring is ∼165–185 K at t =17.9–19.7 yr (Bouchet et al. 2006; Dwek et al. 2008; Seok et al.
2008), while that in SN 1978K is 230 K at t = 29 yr. In addition,
the IR to X-ray flux ratio is of the order of unity in both objects(For SN 1978K, FX ∼ 1 × 10−12 erg cm−2 s−1 ; Smith et al.
2007). The similarity seems to support the shock-heated CS
dust as an origin of IR emission from SN 1978K.
We estimate the mass of the emitting CS dust swept by theforward shock. CS dust is subject to the destruction by the SN
forward shock. The timescale of sputtering is given byng −1
1 da −1a6τsput ≡
10 yr
(4)
a dt
1 μm
1 cm−3(Dwek & Arendt 1992), where ng is the number density of shocked gas. Then, the evolution of the dust radius is estimated as a(t) = a0 (1 − t/τsput), where τsput
=
6
−3 −1
10 yr(a0 /1 μm)(ng /1 cm ) . Hereafter, the subscript “0”denotes the value before the dust destruction. The gas massloss rate and the CS dust density are related as follows:6

The Astrophysical Journal, 749:173 (12pp), 2012 April 20

Tanaka et al.

For small mass-loss rates, the destruction is not so effectivethat the shocked CS dust mass is roughly proportional to the
mass-loss rate. For large mass-loss rates, the surviving dust
mass is not sensitive to the mass-loss rate. Because of thehigh gas density, the destruction is effective enough that themass of dust swept up by the shock compensates for the mass
destroyed by the shock.8 In the calculation above, we assume
there is no dust-free cavity created by the SN outburst. Since
most of surviving dust is located near the SN shock (atR ∼ 5 × 1017 (t/30 yr)(vSN /5000 km s−1 ) cm), which is larger
than the cavity size (∼7 × 1016 cm). Thus, the results presented
here are not largely affected by the presence of the dust-freecavity.
From the IR SED of SN 1978K, we estimate the mass ofemitting dust to be 1.3 × 10−3 M . The model of the shocked
CS dust with Ṁ = 10−4 M yr−1 and vSN = 4000 km s−1 isconsistent with the observed dust mass if the initial dust radius isa0 = 0.3 μm. A similar conclusion has been reached by Dwek
et al. (2008) for SN 1987A.
We also check the expected temperature for the shocked CS
dust. With Ṁ = 10−4 M yr−1 , the shocked gas density is
about 6 × 103 cm−3 . For the electron temperature about 107 K,
the expected dust temperature is Tdust ∼ 270 and 220 K forthe dust radius a = 0.1 and 0.3 μm, respectively (Dwek et al.
2008). Thus, the expected temperature of the shocked CS dustis also in good agreement with the observations. The surviving,shocked CS dust is distributed in a thin shell whose inner
radius is about 95 % of the radius at the forward shock. Using
the mass absorption coefficient of dust κν for astronomicalsilicate at 24 μm (Draine 2003b), the optical depth of dust isabout 4 × 10−4 (κν /562 cm2 g−1 ), which is consistent with ourassumption of optically thin dust (Section 3). In summary, theIR emission from the shocked CS dust can naturally explain theobserved IR emission from SN 1978K.

-2

CS dust mass (M )10

no destruction(a0 = 0.1 μm)-3

SN 1978K

10

10-4
a0 = 0.3 μm
a0 = 0.1 μm
a0 = 0.01 μm
10-5 -6
10

-5

-4

-3

10
10
10
-1
Mass loss rate (M yr )10-2

Figure 5. Mass of shocked CS dust that survives at t = 30 yr after the explosion
as a function of mass-loss rate. The dashed, solid, and dotted lines show themodel with a0 = 0.3, 0.1, and 0.01 μm, respectively. The red dashed line shows
the model without dust destruction with a0 = 0.1 μm. The other parameters are
set as follows: vSN = 4000 km s−1 , vwind = 10 km s−1 , and f = 200.
(A color version of this figure is available in the online journal.)Ṁ = 4π R 2 ρwind vwind = 4π R 2 fρdust,0 vwind , where ρwind is themass density of the wind gas. We also define the mass density of dust before the destruction; ρdust,0 ≡ (4/3)πρa03 ndust,0 ,where ndust,0 is the number density of dust before the destruction. The CS dust density after the destruction is given byρdust (t) = (4/3)π ndust,0 ρa 3 = (4/3)π ndust,0 ρa03 (1 − t/τsput)3 =3ρdust,0 (1 − t/τsput ) .
Using these relations, the mass of the emitting, shocked CS
dust is described as follows: t
MCSdust (t) =4π R 2 ρdust (t  )vSN dt 0
 t=
4π R 2 ρdust,0 vSN (1 − t  /τsput )3 dt 0ṀwindvSN−4= 8 × 10 M
10−5 M yr−1
5000 km s−1
 t




3 
 v
−1 f −1wind
0 (1 − t /τsput ) dt
×.
10 km s−1
200
30 yr

4.2.3. SN Dust

The last possibility for the origin of the IR emission is dust
formed in SN ejecta. Theoretically, SNe are expected to form
dust masses as large as Mdust = 0.1–1.0 M at the inner layer of
the SN ejecta (Kozasa et al. 1989, 1991; Todini & Ferrara 2001;
Nozawa et al. 2003, 2008, 2010; Bianchi & Schneider 2007;
Cherchneff & Dwek 2010). In fact, such a large amount of dust
is required to explain the observed dust mass in high-redshiftquasars, if the dust is formed only by SNe (Dwek et al. 2007;
Meikle et al. 2007; Gall et al. 2011).
Observationally, the mass of dust formed by SNe is quite
controversial. IR observations of young SNe (t  1–2 yr) in
external galaxies usually find only 10−5 –10−3 M of dust (e.g.,
Wooden et al. 1993; Ercolano et al. 2007; Elmhamdi et al.
2003; Meikle et al. 2007, 2011; Kotak et al. 2009; Szalai et al.
2011). On the other hand, observations of SNRs in our Galaxy
and Large/Small Magellanic Clouds detect a larger amount of
dust. For example, by intensive observations of Cas A at IR tosubmillimeter wavelengths, it is estimated that ∼0.1 M of dust
is formed in total (Rho et al. 2008; Sibthorpe et al. 2010; Barlow
et al. 2010). Recently, Matsuura et al. (2011) detected a large
amount of cold dust in SN 1987A at t = 23.3 yr in far-IR and
submillimeter wavelengths. The estimated dust mass is as large
as ∼0.4–0.7 M , and the dust temperature is ∼20 K.

(5)
At the limit of no dust destruction, the integral in the last parentheses becomes t, and the shock-heated CS dust is proportionalto the mass-loss rate. But τsput
is a function of the shocked gas
density, and thus, the mass-loss rate. With a larger mass-loss
rate, the sputtering timescale becomes shorter, and the integral
in the last parentheses becomes smaller than t.
Figure 5 shows the shocked (but surviving) mass of the
CS dust at t = 30 yr as a function of mass-loss rate. The
dashed, solid, and dotted lines in black show the models with
a0 = 0.3, 0.1, and 0.01 μm, respectively. The red solid line is
the model without dust destruction with a0 = 0.1 μm. Other
parameters are assumed as follows: vSN = 4000 km s−1 ,vwind = 10 km s−1 , and f = 200. The SN shock velocity
of SN 1978K is estimated to be vSN < 4000 km s−1 by the very
long baseline interferometry image, which marginally resolves
it (Smith et al. 2007). Thus, we assume vSN = 4000 km s−1 .

8

Fox et al. (2010, 2011) also derived the expression for the shock-heateddust mass, which is independent on the mass-loss rate. Since they assume adestruction-free shell, the mass of CS dust is larger than our calculation by afactor of four.

7

The Astrophysical Journal, 749:173 (12pp), 2012 April 20

Tanaka et al.

To estimate the IR emission from SN dust, we first consider
the total luminosity available, which gives an upper limit of the
IR luminosity. At the transitional phase, one of the major heating
sources of SN dust formed at the inner layers of the ejecta is theradioactive decay of 44 Ti, which gives L ∼ 1036 –1037 erg s−1(Fransson & Kozma 2002). This is much smaller than the IR
luminosity (LIR ∼ 1.5 × 1039 erg s−1 ) of SN 1978K.
A possible additional heating source is UV or X-ray emission
from the ejecta–CSM interacting region. X-ray luminosity of
SN 1978K is very high (3×1039 erg s−1 ; Smith et al. 2007), being
enough to account for the observed IR luminosity. However,the X-ray emission from the shocked region is absorbed by
the optically thick ejecta via photo-absorption even at thetransitional phase (Fransson & Chevalier 1987), and it cannot
reach the inner layer and cannot heat the SN dust there
effectively. Therefore, SN dust at the inner layers of the ejecta
is quite unlikely to be the origin of the IR emission fromSN 1978K.
Note that dust formation in a cool dense shell has been
proposed for some interacting SNe (e.g., Pozzo et al. 2004;
Smith et al. 2008, 2009; Mattila et al. 2008; Fox et al. 2009;
Chugai 2009). The cool dense shell may be formed at thereverse-shocked region, and thus, the dust can be heated by
the UV/X-ray radiation. Although it is not clear whether
1.3 × 10−3 M of the dust can be formed in the cool denseshell, this scenario cannot be completely excluded as a possible
origin for the IR emission from SN 1978K.

Luminosity (erg s-1)

1040
IS dust echo
CS dust echo

SN 1978K

1039

ShockedCS dust1038

1037
SN dust
10

36

0

10

20
30
Time (yr)40

50Figure 6. Expected IR luminosity for normal SNe. The red and blue lines show
the echo of the IS dust and the CS dust, respectively. For the IS dust echo, thered solid and dashed lines show the case for n = 0.1 and 1 cm−3 , respectively.
For the CS dust echo, the solid line shows the case for Ṁ = 10−5 M yr−1 andtwind = 106 yr (outer radius R2 = 3 × 1019 cm). The expected luminosities of
the SN dust and the shocked CS dust are indicated with the arrows.
(A color version of this figure is available in the online journal.)

which has the largest contribution to the total luminosity. Suppose that a dust grain is irradiated by an SN with the luminosity LSN at a distance D. For simplicity, the spectrumof the luminous source is assumed to be blackbody with atemperature Ts . Then, the absorbed luminosity is given by
Labs = (LSN /4π D 2 )π a 2 Qabs (Ts ) , where Qabs (T ) is the
Planck mean of the absorption efficiency factor. When the dust
particle emits radiation with a temperature Tdust , the emitted
4
luminosity is Lrad = 4π a 2 σ Tdust
Qabs (Tdust ) . Under the equilibrium, i.e., Labs = Lrad , we can derive the temperature as
follows:
1/6
−1/3 
LSN
D
Tdust 140 K
1.4 × 1019 cm
6 × 1042 erg s−1
1/6 
−1/6
 abs
a
Q (Ts )×.
(6)
1
0.1 μm

5. IR EMISSION FROM NORMAL SNe
From the upper limits of the flux for five non-detected SNe,upper limits of the emitting dust mass can be obtained. In thissection, we first summarize expected IR emission from normal
SNe at the transitional phase. Then, we discuss the implicationof the non-detection by current observations.
5.1. Summary of IR Emission at the Transitional Phase
For CSM-rich SN 1978K, we have discussed IR emission
expected by (1) light echo from CS and IS dust, (2) shocked
CS dust, and (3) SN dust. In this section, we summarize the IR
emissions from normal SNe by these processes, presenting the
expected IR luminosity and typical dust temperature.
We have shown the light curve of the light echo by CS
dust for Ṁ = 10−4 M yr−1 in Figure 4. Using the same
method, we calculate the echo luminosity with a normal massloss rate Ṁ = 10−5 M yr−1 , and a normal SN luminosity
L0 ∼ 6 × 1042 erg s−1 . The solid blue line in Figure 6 shows theIR luminosity with the wind duration twind = 106 yr. The echo
luminosity is about L = 3 × 1036 erg s−1 at t = 30 yr.
Figure 6 also shows the luminosity by IS dust echo as
calculated in Section 4.2.1. The light curve of the IS dust echo
has a long timescale, and will exceed that of the CS dust echo
after t = 5–10 yr for the ISM density of n = 0.1–1.0 cm−3 .
Note that a typical IS density is not well constrained. Gaustad &van Buren (1993) derived a filling factor of IS gas with density>0.5 cm−3 to be only about 15%. Dring et al. (1996) derived aneven smaller filling factor: 8.5% for the gas density >0.01 cm−3 .
When the density is lower than 0.1 cm−3 , the IR luminosity by
the echo is roughly proportional to the gas density since the IS
dust is optically thin.
For both CS and IS dust cases, the typical temperature of
IR emission can be roughly estimated at the closest pointfrom the SN, i.e., at D = ct/2 = 1.4 × 1019 (t/30 yr) cm,

2
Here, we used Qabs (Tdust ) /a = 0.127Tdust, which is derived
for T  100 K by Dwek (1987). We find that this formula isapplicable up to T ∼ 250 K for our purpose. The expected
temperature for the light echo is ∼140 K at t = 30 yr.
Next we consider the shocked CS dust. Figure 5 shows thatthe typical mass of the surviving, shock-heated dust is about
3 × 10−4 M at t = 30 yr for Ṁ = 10−5 M yr−1 andvSN = 5000 km s−1 . With this mass-loss rate, the shockedgas density is about 6 × 102 cm−3 . For the electron temperature
about 107 K, the expected dust temperature is Tdust ∼ 160 K forthe grain radius of a = 0.1 μm (Dwek et al. 2008). The total
IR luminosity can be derived by integrating Equation (1) overwavelength:
4π 3 −1
4
LIR = Mdust
ρa
4π a 2 σ Tdust
Qabs (Tdust )
3


Mdust
1.8 × 1037
0.1 M
−1 Tdust 6ρerg s−1 ,(7)×
3.8 g cm−3
50 K

where, as in Equation (6), we adopt Qabs (Tdust ) /a =2
0.127Tdust. Using this equation, the total luminosity of theshocked CS dust is estimated as ∼6 × 1037 erg s−1 . For a higher
8

The Astrophysical Journal, 749:173 (12pp), 2012 April 20
1

Table 5Upper limit of Dust Mass

NGC 1313

10

Tdust = 100 K-2
4 x 10 M
2 x 10-2 M

Flux (mJy)

SN 1962M
Amorphous
carbonAstronomicalsilicate

0

10

10-1

10-2

Tanaka et al.

SN

101

SN 1909A
SN 1917A
SN 1951H
SN 1962M
SN 1968D

10
Wavelength (μm)

SN 1909A
SN 1917A
SN 1951H
SN 1962M
SN 1968D

100

Flux (mJy)

Amorphous
carbonAstronomicalsilicate
Tdust = 100 K
4 x 10-1 M-1
2 x 10 M

10-1
Tdust = 200 K
1 x 10-2 M
7 x 10-3 M
1

101

10
Wavelength (μm)

Flux (mJy)

Amorphous
carbonAstronomicalsilicate

Tdust = 200 K-3
4 x 10 M-3
2 x 10 M
1<3 × 10−3<3 × 10−2<4 × 10−3<5 × 10−4<1 × 10−2

10

99
90
56
44
39

<3 × 10−2<3 × 10−1<5 × 10−2<2 × 10−2<2 × 10−1<1 × 10−3<1 × 10−2<2 × 10−3<2 × 10−4<7 × 10−3

Figure 7 shows SEDs of five SNe in the three different
host galaxies. These upper limits are compared with the model
spectra of amorphous carbon (red) and astronomical silicate(blue) with Tdust = 100 and 200 K (solid and dashed lines,
respectively). The assumed dust mass is shown in the panels. In
Table 5, we show the upper limit of the dust mass for the case ofamorphous carbon and astronomical silicate for the temperature
Tdust = 100 and 200 K.
The upper limit of the dust mass strongly depends on the dust
temperature. The blue dashed line in Figure 8 show the upper
limits of the dust mass for SN 1962M as a function of the dust
temperature for the case of astronomical silicate. Constraints
on the other four SNe are weaker than those for SN 1962M
by a factor of 2–20. For comparison, Figure 8 also shows theestimated dust mass and temperature of SN 1978K (this paper)and both warm and cold component of SN 1987A (Bouchet
et al. 2004; Matsuura et al. 2011).
The solid lines in Figure 8 show the dust mass as a functionof dust temperature for LIR = 1036 , 1037 , 1038 , and 1039 erg s−1(from bottom to top). This can be analytically derived from
Equation (7):


LIR
−2
Mdust = 5.6 × 101037 erg s−1
−6 


ρ
Tdust
M .
(8)
×
50 K
3.8 g cm−3

-1

10-2

<7 × 10−2<7 × 10−1<1 × 10−1<4 × 10−2<4 × 10−1

5.2. Constraints from Non-detections

Tdust = 100 K-2
7 x 10 M-2
3 x 10 M

SN 1909A
SN 1951H

10

100

M101

100

99
90
56
44
39

temperature can be Tdust < 50 K. Even if only Mdust = 10−3 Mof dust is formed, the temperature is about 100 K, which islower than the expected temperatures for the light echo and theshocked CS dust.
In summary, IR emission of normal SNe at the transitional
phase can result from three types of mechanisms, i.e., IS dustecho, shocked CS dust, and SN dust. IS dust echo and shocked
CS dust have a relatively high temperature (Tdust = 150–270 K)
while the SN dust has a lower temperature (Tdust < 50–100 K).
The total luminosity of the IS echo can be LIR > 1037 erg s−1
for the IS density of n > 0.1 cm−3 (Figure 6). The shocked CS
dust could also have LIR ∼ 6 × 1037 –4 × 1039 erg s−1 for Ṁ =10−5 –10−4 M yr−1 (see the arrow in Figure 6). The luminosityof the SN dust must be LIR < 1037 erg s−1 .

SN 1917A
SN 1968D

10-2

Mdust ( M )Tdust = 200 K

Astronomical silicate

NGC 6946

100

Mdust ( M )Tdust = 100 K

Amorphous carbon

Tdust = 200 K-4
5 x 10 M-4
2 x 10 M

1

Age
(yr)

100

Wavelength (μm)

Figure 7. Upper limits of the flux for five non-detected SNe in NGC 1313,
NGC 6946, and M101 (from top to bottom). The lines show the model spectraof amorphous carbon (red) and astronomical silicate (blue) with Tdust = 200 K(dashed) and 100 K (solid). The dust mass assumed in the model is shown in
the panels.
(A color version of this figure is available in the online journal.)

mass-loss rate Ṁ = 10−4 M yr−1 , the shocked CS dust
mass becomes Mdust ∼ 1 × 10−3 M , and the temperature is
Tdust ∼ 270 K. The total IR luminosity becomes 4×1039 erg s−1 .
Finally, let us consider the SN dust. For the SN dust, we
assume LIR < 1037 erg s−1 , which is expected by the radioactive
decay of 44 Ti. If the mass of SN dust is Mdust = 0.1 M , theFor these lines in Figure 8, we do not use the analytic formulae,
but numerically integrate Equation (1) over wavelength.
9The Astrophysical Journal, 749:173 (12pp), 2012 April 20

Tanaka et al.

1

10

1

Observational limit for SN 1962M(astronomical silicate)010Mdust (M )

0

10

87A
cold

Flux (mJy)

-1

10

10

1039
-2

10

-3

10

10
1036

38

1037

Herschel

ALMA
SPICA

10-2
JWST

78K
-3

87A
warm

10-4

Confusion limit (3m)

10-1

-1

L (erg s )SN at 5 Mpc (t = 30 yr)
Astronomical silicate
SN dustShocked CS dust
IS dust echo

10

ShockedCS dust

-4

10

-5

10

0

50

100

150
Tdust (K)

200

250

300

10

100
Wavelength (μm)

1000

Figure 9. Expected IR SEDs of SN dust (purple), shocked CS dust (blue), andIS dust (red) compared with 5σ detection limit of 1 hr imaging observation
with various telescopes, i.e., JWST (black), SPICA (red), Herschel (black),and ALMA (black). Gray dashed line shows the expected confusion limitof background galaxies with a 3 m telescope (T. T. Takeuchi et al. 2012, in
preparation). The distance to the SN is set to be 5 Mpc. For the SN dust,
we assume Mdust = 0.1 M with Tdust = 50 K. This roughly corresponds tothe total IR luminosity of LIR ∼ 1037 erg s−1 . For the shocked CS dust, we
assume Mdust = 3 × 10−4 M with Tdust = 150 K (the total luminosity is
LIR = 3.9 × 1037 erg s−1 ). For the IS dust echo, we assume LIR ∼ 1037 erg s−1
with Tdust = 150 K. The effective emitting mass of dust is Mdust ∼ 10−4 M .
(A color version of this figure is available in the online journal.)

Figure 8. Upper limit of the dust mass as a function of the dust temperature
for the case of astronomical silicate. The dashed blue line shows the upper
limit for SN 1962M, which is derived from the non-detection in our current
observation. Solid curves show the upper limit for a given total luminosity
L = 1036 –1039 erg s−1 . Our current observations are sensitive to the luminosity
L  1038 erg s−1 . The observations of SNe 1978K (pink diamond) and1987A (black box; Bouchet et al. 2004; Matsuura et al. 2011) are also
shown for comparison. The expected ranges for the shocked CS dust withṀ = 10−5 M yr−1 (lower left) and Ṁ = 10−4 M yr−1 (upper right) are
shown (for dust size a = 0.1–0.3 μm).
(A color version of this figure is available in the online journal.)

The figure shows that future observations at <30 μm aresensitive to the shocked CS dust or IS dust echo. These twocases may show similar overall spectra, but ionic emission
lines are also expected for the shocked CS dust. Low-resolution
spectroscopy will be possible with SPICA, which will be useful
to distinguish the shocked CS dust emission from IS dust echo.
IR emission by SN dust has a smaller contribution to the IR
luminosity. However, thanks to the lower temperature than theshocked CS dust and IS dust echo, the expected spectrum has apeak at longer wavelengths. The contamination by the shocked
CS dust may be small if Ṁ  10−5 M yr−1 . In addition, theemission from the IS dust echo can also be small if a typical IS
density is less than 0.1 cm−3 (Gaustad & van Buren 1993; Dring
et al. 1996). If the contributions from these two components are
small, the detection of SN dust is limited only by the confusion
limit of background galaxies, which is an order of 0.1 mJy at
50 μm with a 3 m telescope.
The great advantage to detect the SN dust at the transitional
phase is the low optical depth. If an SN forms 0.1 M of dust,
the optical depth becomes



Mdust
3
κν
−2
κν Mdust R
0.1
τν =
4π
562 cm2 g−1
0.1 M

−2


−2
t
vSN
.
(9)
×
5000 km s−1
30 yr

As shown in Figure 8, our current observations are sensitive
only to the IR luminosity as high as >1038 erg s−1 . This limiting
luminosity is much higher than the predictions for the SN dust.
The luminosity of the IS dust echo can be 1 × 1038 erg s−1
with 1 cm−3 . Our current observations can only exclude ahomogeneous IS density of n ∼ 10 cm−3 .
For the shocked CS dust, the expected ranges of the dust
mass and temperature are shown by two gray boxes in Figure 8.
The left lower box shows the range for Ṁ = 10−5 M yr−1 ,while the right top for Ṁ = 10−4 M yr−1 , Our current best
observational limit excludes the existence of the shocked CS
dust with Ṁ = 10−4 M yr−1 .
6. FUTURE PROSPECTS
We show that our current observations are sensitive tothe luminosity L > 1038 erg s−1 . In this section, we discussprospects for future observations. Figure 9 shows the expected
flux for the SN dust (purple, Mdust = 0.1 M and Tdust = 50 K),
the shock-heated CS dust (blue, Mdust = 3 × 10−4 M , andTdust = 150 K, expected for Ṁ = 10−5 M yr−1 ), and the IS
dust echo (red, L = 1037 erg s−1 and Tdust = 150 K, the effective
emitting dust mass is Mdust ∼ 10−4 M ). The SN is assumed to
be located at 5 Mpc.
The model spectra are compared with 5σ detection limit of
1 hr imaging observations with various telescopes, i.e., with The
James Webb Space Telescope (JWST, black; Wright et al. 2004;
Swinyard et al. 2004), SPICA (red; Nakagawa 2010; Wada &Kataza 2010; Swinyard 2008; Kataza et al. 2010), Herschel(black; Pilbratt et al. 2010; Poglitsch et al. 2010; Griffin et al.
2010), and ALMA (black; Wootten & Thompson 2009). For
SPICA, we also show the limit for spectroscopic observations(R = 50) in the red dashed line (Kataza et al. 2010). In the
Appendix, we show similar plots for other dust species.

For the mass absorption coefficient of dust κν , we take thevalue for astronomical silicate at 24 μm (Draine 2003b). At
longer wavelengths, the coefficient is lower (κν ∝ λ−2 ). At
early epochs, such a large mass of dust is optically thick even at
IR wavelengths. On the contrary, at the transitional phase, theoptical depth can be smaller than unity, which enables a reliable
estimate of the total dust mass.
10

The Astrophysical Journal, 749:173 (12pp), 2012 April 20

Tanaka et al.

Another advantage is the small confusion with the IS dust
swept up by the SN ejecta. The mass of the IS dust (MISdust )can be approximately written as follows by neglecting thedestruction:
4π
MISdust =(vSN t)3 nISM mp f −13
3  t 3vSN
2 × 10−65000 km s−1
30 yr  f −1
 n
ISM
×
M .
(10)
1 cm−3200

101
0Flux (mJy)

10

SN at 5 Mpc (t = 30 yr)
Forsterite
SN dustShocked CS dust
IS dust echo

Confusion limit (3m)
Herschel

ALMA

-1

10

SPICA
-2

10

JWST
10-3
-4

The mass of swept-up IS dust is nearly 0.1 M in SNRs at
t = 1000 yr, which makes it difficult to estimate the SN dust
mass in SNRs. On the contrary, the swept-up IS dust mass at thetransitional phase is only an order of 10−6 M .
In summary, future observations of SNe at the transitional
phase will fill the gap of infrared observations of SNe with theage of 10–100 yr. At <30 μm, observations will provide an interesting way to study the CS and IS environments of the progenitor. With SPICA, these components can also be studied by lowresolution spectroscopy. If these two components are relatively
weak, the SN dust might be detected at longer wavelengths,
which enables the reliable measurement of the dust mass.

10

10

1

10

0

Flux (mJy)

10100
Wavelength (μm)

SN at 5 Mpc (t = 30 yr)
Amorphous carbon
SN dustShocked CS dust
IS dust echo

1000

Confusion limit (3m)
Herschel

ALMA

-1

10

SPICA
-2

10

JWST
10-3

7. CONCLUSIONS
SN explosions have been studied mostly by observationsof extragalactic SNe at the early phase (t < a few years)
or observations of SNRs at the late phase (t > 100 yr).
Observations at the transitional phase from SN to SNR havenot been extensively performed, especially at IR wavelengths.
We present theoretical predictions for the IR emission from SNeat the transitional phase. We show that the emission arises fromSN dust, light echo by CS and IS dust, and shocked CS dust.
We search for IR emission toward six core-collapse SNe at thetransitional phase in the nearby galaxies NGC 1313, NGC 6946,
and M101 by using the data taken with the AKARI satellite
and Spitzer. Among the targets, we detect emission associatedwith SN 1978K in NGC 1313. The emission is explained by
1.3 × 10−3 M of silicate dust. The shocked CS dust for gas
mass-loss rate Ṁ = 10−4 M yr−1 is the most probable originof the IR emission from SN 1978K. Dust formed in the outer,cool dense shell might also be a possible origin.
IR emission from the other five objects is not detected.
Current observations are sensitive only to the total luminosityof L > 1038 erg s−1 . The non-detection of SN 1962M excludes
the existence of the shocked CS dust for the high gas mass-loss
rate Ṁ = 10−4 M yr−1 .
Future observations will fill the gap of the IR observations at
the transitional phase. At <30 μm, the emission is dominatedby the shocked CS dust and IS dust echo. The observations
will provide an interesting opportunity to study the CS and IS
environments of the SN progenitor. SN dust component has
a smaller contribution, but has a spectrum peaking at longer
wavelengths. If the emission by the shocked CS dust and IS
dust echo does not dominate the IR emission, the SN dust might
be detected at >30 μm. This will give a reliable estimate of the
SN dust mass, which does not suffer from the high optical depthof dust and the confusion with the IS dust.

-4

10

10

100

1000

Wavelength (μm)

1010Flux (mJy)

10

SN at 5 Mpc (t = 30 yr)
Graphite
SN dustShocked CS dust
IS dust echo

Confusion limit (3m)
Herschel

ALMA

-1

10

SPICA
-2

10

JWST
-3

10

-4

10

10

100

1000

Wavelength (μm)

Figure 10. Expected SEDs of SNe at the transitional phase (t = 30 yr) for
forsterite, amorphous carbon, and graphite (from top to bottom). For the shocked
CS dust, Tdust = 150 K and Mdust = 3 × 10−4 M are assumed. For the IS
dust echo, Tdust = 150 K and Mdust = 1 × 10−4 M are assumed, whichgive the luminosity of L = 1037 erg s−1 . For the SN dust, Tdust = 50 K and
Mdust = 0.1 M are assumed.
(A color version of this figure is available in the online journal.)

comments, and G. Folatelli for valuable discussions. This research is based on observations with AKARI, a JAXA projectwith the participation of ESA, and observations made with the
Spitzer Space Telescope, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under a contract with NASA. We have made use of data products from theTwo Micron All Sky Survey, which is a joint project of the
University of Massachusetts and the Infrared Processing andAnalysis Center/California Institute of Technology, funded by
the National Aeronautics and Space Administration and the National Science Foundation. This research has been supported in

We are grateful to the referee for giving useful comments.
We thank T. T. Takeuchi for providing the model of the confusion limit of SPICA. We also thank T. Kozasa for useful11The Astrophysical Journal, 749:173 (12pp), 2012 April 20

Tanaka et al.

part by World Premier International Research Center Initiative,
MEXT, Japan, and by the Grant-in-Aid for Scientific Researchof the Japan Society for the Promotion of Science (20340038,
22684004, 22840009).

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APPENDIXSEDs FOR AMORPHOUS CARBON, GRAPHITE,AND FORSTERITEAs discussed in Section 6, future mid/far-IR observations
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