Rironkon symposium 2019, NAOJ, 12/25-27

星周円盤と相互作用する超新星の2次元放射流体シミュレーション
Suzuki, Moriya, & Takiwaki (2019), ApJ accepted (arXiv:1911.09261)

鈴木昭宏 (国立天文台科学研究部, NAOJ fellow)

collaborator: 守屋尭, 滝脇知也 (国立天文台科学研究部)

Outline• Motivation• Method• Simulations and Results• SummarySN 1987A

1984

1987

重力崩壊型超新星
Core collapse•

光分解による鉄コアの重力崩壊
Fe•

衝撃波形成と伝搬•

超新星ショックブレイクアウト•

超新星エジェクタの膨張と熱放射Core bounce

⁵⁶Fe+γ→13⁴He+4n
⁴He+γ→2p+2np+e-→ν+ne

Shock propagation

Luminosity

50 dρc =3 10¹⁴[g/c.c.]Shock breakout

100 d1042-1043erg/s

SN explosion

UV/X-ray ﬂash
post shock 〜0.1keV

Days since detection
Light curves of various SNe (Filippenko 1997)

Optical重力崩壊型超新星
SN spectranoyesHydrogenno

Modjaz+(2014)

Type I

Type II(+Type IIn)

yes

SiliconnoyesType Ia (+Type Ia-CSM)
thermonuclear SNe

HeliumType Ic

Type Ib stripped envelope SNe(+Type Ibn)
+broad-lined Ic (Ic-BL)

back into the SN ejecta. In the simplest picture, four zones are delineated where
heated material can radiate and contribute to the observed spectrum:

Interacting SNe: type IIn, Ibn, and Ia-CSM

Zone 1. The unshocked CSM outside the forward shock (photoionized).
Zone 2. The swept-up CSM that has been hit by the forward shock.
Zone 3. The decelerated SN ejecta that have encountered the reverse shock.
Zone
4. The freely expanding
SN ejecta. (CSM:星周物質)
Circumstellar
matterIn normal SNe, we only see emission from Zone 4 at visible wavelengths as
the photosphere recedes back through the freely expanding ejecta, creating broad
emission and absorption lines whose velocities decrease smoothly with time as theFig. 1 Sketch of the basic picture in CSM interacting SNe. Four different zones are noted withnumbers: (1) the pre-shock CSM, (2) the shocked CSM, (3) the shocked SN ejecta, and (4) the
freely expanding SN ejecta. These zones are divided by boundaries corresponding to the forward
shock, the reverse shock, and the contact discontinuity between the shocked CSM and shocked
from
Smith
(2016)
ejecta where material cools, mixes via Rayleigh-Taylor instabilities, and piles
up. This
is oftentype IIn SN 2006gy
Smith+(2010)

back into the SN ejecta. In the simplest picture, four zones are delineated where
heated material can radiate and contribute to the observed spectrum:

Interacting SNe: type IIn, Ibn, and Ia-CSM

Zone 1. The unshocked CSM outside the forward shock (photoionized).
Zone 2. The swept-up CSM that has been hit by the forward shock.
Zone 3. The decelerated SN ejecta that have encountered the reverse shock.
Zone
4. The freely expanding
SN ejecta. (CSM:星周物質)
Circumstellar
matterIn normal SNe, we only see emission from Zone 4 at visible wavelengths as
the photosphere recedes back through the freely expanding ejecta, creating broaddissipationoftheejectaenergyemissionand absorption
lines
whose
velocities kinetic
decrease smoothlywith time as the→bright thermal emission

Fig. 1 Sketch of the basic picture in CSM interacting SNe. Four different zones are noted withnumbers:(1) thepre-shock CSM, (2)emission
the shocked CSM,
(3) the松岡ポスター
shocked SN ejecta, and (4) theradio
synchrotron→ P32
freely expanding SN ejecta. These zones are divided by boundaries corresponding to the forward
shock, the reverse shock, and the contact discontinuity between the shocked CSM and shocked
from
Smith
(2016)
ejecta where material cools, mixes via Rayleigh-Taylor instabilities, and piles
up. This
is oftentype IIn SN 2006gy
Smith+(2010)

back into the SN ejecta. In the simplest picture, four zones are delineated where
heated material can radiate and contribute to the observed spectrum:

Interacting SNe: type IIn, Ibn, and Ia-CSM

Zone 1. The unshocked CSM outside the forward shock (photoionized).
Zone 2. The swept-up CSM that has been hit by the forward shock.
Zone 3. The decelerated SN ejecta that have encountered the reverse shock.
Zone
4. The freely expanding
SN ejecta. (CSM:星周物質)
Circumstellar
matterIn normal SNe, we only see emission from Zone 4 at visible wavelengths as
the photosphere recedes back through the freely expanding ejecta, creating broaddissipationoftheejectaenergyemissionand absorption
lines
whose
velocities kinetic
decrease smoothlywith time as the→bright thermal emission

Smith+ (2007)

Fig. 1 Sketch of the basic picture in CSM interacting SNe. Four different zones are noted withnumbers:(1) thepre-shock CSM, (2)emission
the shocked CSM,
(3) the松岡ポスター
shocked SN ejecta, and (4) theradio
synchrotron→ P32
freely expanding SN ejecta. These zones are divided by boundaries corresponding to the forward
shock, the reverse shock, and the contact discontinuity between the shocked CSM and shocked
from
Smith
(2016)
ejecta where material cools, mixes via Rayleigh-Taylor instabilities, and piles
up. This
is oftentype IIn SN 2006gy
Smith+(2010)

back into the SN ejecta. In the simplest picture, four zones are delineated where
heated material can radiate and contribute to the observed spectrum:

Interacting SNe: type IIn, Ibn, and Ia-CSM

Zone 1. The unshocked CSM outside the forward shock (photoionized).
Zone 2. The swept-up CSM that has been hit by the forward shock.
Zone 3. The decelerated SN ejecta that have encountered the reverse shock.
Zone
4. The freely expanding
SN ejecta. (CSM:星周物質)
Circumstellar
matterIn normal SNe, we only see emission from Zone 4 at visible wavelengths as
the photosphere recedes back through the freely expanding ejecta, creating broaddissipationoftheejectaenergyemissionand absorption
lines
whose
velocities kinetic
decrease smoothlywith time as the→bright thermal emission

濃い星周物質はどこからきたのか? Smith+ (2007)

Fig. 1 Sketch of the basic picture in CSM interacting SNe. Four different zones are noted withnumbers:(1) thepre-shock CSM, (2)emission
the shocked CSM,
(3) the松岡ポスター
shocked SN ejecta, and (4) theradio
synchrotron→ P32
freely expanding SN ejecta. These zones are divided by boundaries corresponding to the forward
shock, the reverse shock, and the contact discontinuity between the shocked CSM and shocked
from
Smith
(2016)
ejecta where material cools, mixes via Rayleigh-Taylor instabilities, and piles
up. This
is oftentype IIn SN 2006gy
Smith+(2010)

Summary of Progenitor and Explosion Properties of Each of the Different Models
Modelalog Teffb(K)

log Lc
(Le)

Rphd
(Re)

Route
(Re)

Ebin,envf(1047 erg)

3.52
3.53
3.54
3.54
3.52
3.51
3.59

4.93
6.18
6.15
6.17
4.96
5.23
5.89

876.3
3553.23282.73363.2906.4
1316.7
1899.7926.95571.03658.46081.2960.6
1434.78977.9

8.67
2.14
0.91
3.10
8.43
5.94
1.25

大質量星からの質量放出
No injection
Ldep5d39uni
Ldep5d39base
Ldep5d39middle
Ldep1d38uni
Ldep1d39uni
Ldep1d40uni• 重力崩壊直前(10-100yrs)に0.1- 1 M◉/yr
Notes.

Unbound Massg
(Me)

tSBh
(day)

Lpt,50i(10 erg s−1)42Plateau Durationj
(day)

0.00
1.68
2.35
89.1
7.14
9.02
9.04
98.7
7.36
5.49
7.23
114.6
6.30
9.07
8.96
113.6
HST
image
(© 87.9
NASA)
0.00
1.75 of η Carinae
2.34
0.00
2.56
3.58
84.6
6.73
12.88
12.33
105.0

a

のmass-loss?

Model name.
Effective temperature of the progenitor at the time of core collapse.
cPhotospheric luminosity of the progenitor at the time of core collapse.
dPhotospheric radius of the progenitor at the time of core collapse.
e
The radial coordinate of the outermost numerical cell of the progenitor at the time of core collapse.
bfBinding energy of the envelope of the progenitor at the time of core collapse. This is calculated as Ebin,env = ò

MtotalMHe

Gm /r dm, where Mtotal and MHe are the total

mass and He core mass of the progenitor, respectively.
g
The integrated mass of the progenitor at the time of core collapse for the region in the envelope where the speciﬁc total energy is positive. Here, the total energy is
the sum of the kinetic energy, internal energy, and the gravitational potential energy.
hTime since the explosion to the shock breakout.
iLuminosity of the supernova at 50 days since the shock breakout.
jDuration of plateau. This is calculated as the time interval between the time of shock breakout and the time when the luminosity becomes half of Lpt,50.

• LBV-like giant eruption?

• wave-driven mass-loss?
0.24 cm gfor a pure metal composition (Z=1). We
2−1

employ the analytic equation of state given by Paczynski

Quataert&Shiode (1983),(2012),(2017),from…radiation, ions, andwhich Fuller
contains contributions
electrons. We trace the ionization fractions of hydrogen andhelium solving the Saha equations.
The code does not include◉ a nuclear reaction network, and56
Ni is given by hand. Here, the mass of Ni is assumed to be
Woosley (2012), M…
Ni=0.07 Me, which is a characteristic value for observed
SNe II (Hamuy 2003; Smartt 2009; Dall’Ora et al. 2014;
Müller et al. 2017). Then, it is distributed from the innerboundary up to the mass coordinate of m(r)=3 Me.

• 対生成不安定 (M~100-300M )?
• 連星進化による質量放出?

3. Result
In Table 1, we summarize the progenitor and explosion
properties for each of the different models. Here, the model
name of Ldep5d39middle, for example, denotes the model with
Ldep=5.0×1039 erg s−1, and MIDDLE distribution. For all
the models, the He core mass and total mass of the progenitors
at the time of core collapse are 4.56 Me and 11.9 Me,
respectively.
We ﬁrst focus on the model with an energy injection rate of
Ldep=5.0×1039 erg s−1 with a UNIFORM distribution(hereafter, Ldep5d39uni). This is to demonstrate how thesustained energy deposition in the envelope, continuing for a
few years before the core collapse, affects the density proﬁle ofthe progenitors and SNe.
Ouchi&Maeda (2019)• 共通外層進化 → 爆発?

Figure 1. Top: the density proﬁle at the time of core collapse. The model with

unity for a are expected regardless of the initial spinstate. In what follows, we will assume that the compact object accretes the entire core mass (facc = 1.0)
and we explore explosion energies between 3 ⇥ 1050 erg
and 1052 erg. We further assume that regardless of the
precise mechanism or energetics, the energy is sharedroughly spherically with the hydrogen envelope (Chevalier 2012; Soker & Gilkis 2018; Soker 2019).

大質量星からの質量放出

3. CIRCUMSTELLAR MATERIAL• 重力崩壊直前(10-100yrs)に0.1- 1EXPELLED
M◉/yr

DURING BINARY COALESCENCE

As a basis for our analysis, we model binary coalescence and the circumbinary ejecta that results from themerger of two example mass ratio binaries. Here we
describe our numerical method, the unstable mass exchange that leads the binary to merge, and the CSM
mass distribution that this runaway mass transfer creates. The distribution of CSM at large scales, r ⇠1015 cm, is of particular importance for the outburst
light curves. We therefore focus our numerical modelingon the early phases of runaway, unstable mass exchangethat expels this largest-scale CSM.

のmass-loss?

• LBV-like giant eruption?
• wave-driven mass-loss?

Quataert&Shiode(2012), Fuller (2017), …
3.1. Hydrodynamic Models of Binary Coalescence

Our models are simulated within the Eulerian hydrodynamic code Athena++ (Stone,◉J.M., in preparation),
Figure 1. Runaway Roche lobe overflow
leading to binary
Schroeder+(2019)
and are based on the methodology described in MacLeod
6
coalescence for binary mass ratio q = 0.1. The upper panel
Woosley (2012),…
et al. (2018a,b,c).
We use a spherical polar coordinateshows binary separation as a function of time during oursystem surrounding the donor star in a binary pair. We
simulation. The lower panel shows four snapshots (marked
model the interaction of this donor star with a softenedwith vertical lines in the upper panel) of the gas density dispoint mass representing an unresolved companion obtribution in the orbital plane. We show these slices throughthe orbital plane in a rotated x0 y 0 coordinate system soject. The domain extends over the full three-dimensionalthat the companion star lies along the +x-axis. The dimensolid angle from 0.1 to 100 times the donor-star’s origisionless simulations are rescaled to donor star mass of 30M
nal radius.
and radius of 1000R for specificity in these images. FollowThe donor star is modeled by polytropic envelope,ing Roche lobe overflow, mass is pulled from the donor starwith structural index = 1.35. The donor has a coreand expelled from the binary system dragging the binary tomass of 25% its total mass. The gas in the simulationtighter separations.
domain follows an ideal-gas equation of state, with index = 1.35. These choices are intended to approximately represent a convective, isentropic envelope of a
The calculations themselves are carried out in dimenmassive star in which radiation pressure is important insionless units where the donor’s mass, radius, and gravFigure 2. Three dimensional distribution of ejecta near the time of merger. As in Figure 1, we have scaled to a fiducial donor
the equation of state (in which case = ! 4/3,
e.g.
constantareall settomerges
unity.
may,
mass
of 30M anditational
radius of 1000R. When thecompact
objectwith They
the donor’score, thereit is surrounded by an extensive,
thick
torus
of
debrisexpelledbythemergeritself.
MacLeod et al. 2017a; Murguia-Berthier et al. 2017).
fore, be rescaled to a physical binary of any mass or size.
We initialize the calculation at a separation slightly
reportdensity
on twomodels whichhavesecondary
to of stellar mass loss
In Figure 2, weBelow,
show the welarge-scale
distribulocities are
similarto other sources• 対生成不安定 (M~100-300M )?
• 連星進化による質量放出?
• 共通外層進化 → 爆発?

非球対称な星周物質?

1170

J. Mauerhan et al.

• 観測的示唆.

• disk-like? bipolar?

Kokubo+ (2019)
KISS15sHα lineFigure 2. Total flux, fractional de-biased polarization Pdb , and position
angle θ for SN 2009ip during its 2012 outburst, at epochs between Sep. 21
and Dec. 6. The three Kuiper/SPOL epochs (red) have been averaged forclarity. The Sep. 21 and 24 MMT/SPOL data (light blue) have been averaged
and binned to 300 Å; all other epochs are binned to 50 Å. Note the ∼90◦
offset in θ between the 2012a phase and the peak of 2012b.

Figure 3. Upper panel: temporal evolution of the total V-band (5050–
5950 Å) polarization for SN 2009ip during the 2012a and 2012b phases.
For reference, an arbitrarily scaled version of the V-band light curve isplotted in the background. Lower panel: evolution in the Q–U plane. ISP isconstrained to <0.2 per cent, illustrated by the black dot near the origin. The
dashed lines represent the position angles θ V = 166◦ and 72◦ measured at
the first epoch on Sep. 21 and peak polarization on Oct. 14, respectively. Ourlatest continuum measurement (5200–5500 Å) for Dec. 6 is marked with a
red cross near the origin, consistent with our ISP limit. The approximate
point-reflection symmetry between 2012a and the peak of 2012b suggests
two separate and roughly orthogonal components of polarization on the sky.
Note the temporary shift in θ associated with the bump in the light curve at∼30 d past peak.

2012 event of SN imposter 2009ip Mauerhan+ (2014)

The observed narrow emission lines are probably dominatedby the emission from H II regions in the host galaxy, which
may not be directly related to KISS15s. Nevertheless, thepresent observations do not rule out the possibility of a narrow
emission line ﬂux contribution from the unshocked CSM
photoionized by the SN emission.
Apparently, the proposed toroidal CSM distribution is not aunique solution. Another possible scenario to explain theblueshifted Hα intermediate-width emission line component is
to consider the (one-sided) jet structure of SN ejecta. To
lines (e.g., Graham et al. 2014; Smith et al. 2017). The ﬂux excess
The host galaxy of KISS15s, SDSSJ030831.67-005008.6,
the Survey,
triple-peakHαemissionline (composed ofin the wavelength range of l rest < 5700 Å generally observed in
was imaged by accountthe SDSS for
Legacyand442,thecentral
MNRAS
1166–1180(2014)−1
blueshifted andcentral emission line and ∼10,000 km s
20
redshifted Hα emission lines) observed in SN2010jp, Smithet al. (2012a) proposed a model in which a two-sided
supernova jet, inclined relative to the observer, passes throughinitially spherically symmetric CSMs and produces two
Figure 10. Illustration of a possible configuration for the 2012b phase, afshocked regions
withblueand SN/CSM
redshiftedvelocities.
Figure 9.
Illustrationof a potential
configuration
during the 2012a It should
Figure 24. Schematic drawing of the proposed ejecta-CSM interaction region
event (not to scale), before the onset of intense CSM interaction. The pres-ter the onset of intense CSM interaction. A toroidal distribution of dense

Downloaded from https://academic.oup.com/mnras/article

polarization vectors from the 2012a and 2012b components, as thelatter rises in luminosity.
Kokubo et al.
By Oct. 11, the source has risen to peak flux. The spectrum hasretained the overall appearance of Oct. 5, while the continuum becomes less blue and broad-line P-Cygni absorption becomes apparSpectropolarimetry of SN 2009ip
1177
Figure 20. Same as Figure 19, but one additional Gaussian is included to account for the blueshifted excess emission component. For the
LISS spectrum, theent for He I/NaD. Our Lick/Kast and Kuiper/SPOL measurements
additional Gaussian component is forced to have the same velocity width as the intermediate-width line (assuming they both are unresolved).
on Oct. 14 show that P has risen to a maximum of ∼1.7 per centand θ has reached ∼72◦ . At this time, strong deficits in polarization
are seen for the intermediate-width emission components of Hα(e.g., Jencson et al. 2016). The strong, broad/intermediate Ca II
andspectra of SNe IIn, which is composed of aand,blend
of extent,
Fe II Hβ.
to a lesser
This is likely to be the result of dilution
fromthestrong,intrinsicallytriplet emission with no associated P Cygni absorption featureseveral other emission lines (Stathakis & Sadler 1991; Stritzingerunpolarised emission lines in the outer
2009ip occupies a compact portion of the Q–U plane, in quadrant II.
Meanwhile, a slight increase in P is seen at 5600–5800 Å,
must be emitted by the same region as that with the strong Hαet al. 2012), is not observed or is signiﬁcantly CSM.
weakerin KISS15s
The points form a slightly elongated cluster that shows a continuous
suggesting a potential association with broad He I λ5876 or NaD
emission line (i.e., the CDS and/or photoionized ejecta; Chugai
compared to SN1988Z (Figure 8).
P-Cygni lines.
trend in wavelength, with the bluest points exhibiting the lowest poFig. 4 (top panels) shows the Stokes spectra and the data distribusition angle in the Q–U plane. The cluster as a whole exhibits a small& Danziger 1994; Dessart et al. 2016).
tion in the Q–U plane for the Oct. 14 Lick epoch, binned to 100 Å.
but non-negligible spread in position angle, which corresponds to
The origin of the broad bump at the wavelengths of
4. Discussion
For reference, we also show scaled logarithmic flux spectra in thea range of ∼10◦ on the sky. The dilution in polarization associatedwith intermediate-width Hα emission is apparent as a protrusion
same figure, which has been normalized with a low-order spline
λrest∼7300 Å in KISS15s spectra is unclear, but it may be due
4.1. Host Galaxy Properties
from the cluster of data points which roughly points towards the(avoiding the strongest emission features for the spline nodes). On
to a blend of broad lines of He I and Ca II and several forbidden
origin. Note that our binning causes the line to be unresolved, andOct. 14, near peak brightness, the entire optical spectrum of SN

The Astrophysical Journal, 872:135 (34pp), 2019 February 20

Downloaded from https://academic.oup.com/mnras/article-abstract/442/2/1166/980021 by National Astronomical Observatory of Japan user on 22 October 2018

• 偏光, 非対称な輝線プロファイル (e.g., Hα, Hβ)非球対称な星周物質?

1170

J. Mauerhan et al.

• 観測的示唆.

• disk-like? bipolar?

Kokubo+ (2019)
KISS15sHα lineFigure 2. Total flux, fractional de-biased polarization Pdb , and position
angle θ for SN 2009ip during its 2012 outburst, at epochs between Sep. 21
and Dec. 6. The three Kuiper/SPOL epochs (red) have been averaged forclarity. The Sep. 21 and 24 MMT/SPOL data (light blue) have been averaged
and binned to 300 Å; all other epochs are binned to 50 Å. Note the ∼90◦
offset in θ between the 2012a phase and the peak of 2012b.

非球対称な星周物質と衝突する超新星の見え方は?
2012 event of SN imposter 2009ip Mauerhan+ (2014)(光度曲線, スペクトル, 偏光……)

Figure 3. Upper panel: temporal evolution of the total V-band (5050–
5950 Å) polarization for SN 2009ip during the 2012a and 2012b phases.
For reference, an arbitrarily scaled version of the V-band light curve isplotted in the background. Lower panel: evolution in the Q–U plane. ISP isconstrained to <0.2 per cent, illustrated by the black dot near the origin. The
dashed lines represent the position angles θ V = 166◦ and 72◦ measured at
the first epoch on Sep. 21 and peak polarization on Oct. 14, respectively. Ourlatest continuum measurement (5200–5500 Å) for Dec. 6 is marked with a
red cross near the origin, consistent with our ISP limit. The approximate
point-reflection symmetry between 2012a and the peak of 2012b suggests
two separate and roughly orthogonal components of polarization on the sky.
Note the temporary shift in θ associated with the bump in the light curve at∼30 d past peak.

The observed narrow emission lines are probably dominatedby the emission from H II regions in the host galaxy, which
may not be directly related to KISS15s. Nevertheless, thepresent observations do not rule out the possibility of a narrow
emission line ﬂux contribution from the unshocked CSM
photoionized by the SN emission.
Apparently, the proposed toroidal CSM distribution is not aunique solution. Another possible scenario to explain theblueshifted Hα intermediate-width emission line component is
to consider the (one-sided) jet structure of SN ejecta. To
lines (e.g., Graham et al. 2014; Smith et al. 2017). The ﬂux excess
The host galaxy of KISS15s, SDSSJ030831.67-005008.6,
the Survey,
triple-peakHαemissionline (composed ofin the wavelength range of l rest < 5700 Å generally observed in
was imaged by accountthe SDSS for
Legacyand442,thecentral
MNRAS
1166–1180(2014)−1
blueshifted andcentral emission line and ∼10,000 km s
20
redshifted Hα emission lines) observed in SN2010jp, Smithet al. (2012a) proposed a model in which a two-sided
supernova jet, inclined relative to the observer, passes throughinitially spherically symmetric CSMs and produces two
Figure 10. Illustration of a possible configuration for the 2012b phase, afshocked regions
withblueand SN/CSM
redshiftedvelocities.
Figure 9.
Illustrationof a potential
configuration
during the 2012a It should
Figure 24. Schematic drawing of the proposed ejecta-CSM interaction region
event (not to scale), before the onset of intense CSM interaction. The pres-ter the onset of intense CSM interaction. A toroidal distribution of dense

Downloaded from https://academic.oup.com/mnras/article

polarization vectors from the 2012a and 2012b components, as thelatter rises in luminosity.
Kokubo et al.
By Oct. 11, the source has risen to peak flux. The spectrum hasretained the overall appearance of Oct. 5, while the continuum becomes less blue and broad-line P-Cygni absorption becomes apparSpectropolarimetry of SN 2009ip
1177
Figure 20. Same as Figure 19, but one additional Gaussian is included to account for the blueshifted excess emission component. For the
LISS spectrum, theent for He I/NaD. Our Lick/Kast and Kuiper/SPOL measurements
additional Gaussian component is forced to have the same velocity width as the intermediate-width line (assuming they both are unresolved).
on Oct. 14 show that P has risen to a maximum of ∼1.7 per centand θ has reached ∼72◦ . At this time, strong deficits in polarization
are seen for the intermediate-width emission components of Hα(e.g., Jencson et al. 2016). The strong, broad/intermediate Ca II
andspectra of SNe IIn, which is composed of aand,blend
of extent,
Fe II Hβ.
to a lesser
This is likely to be the result of dilution
fromthestrong,intrinsicallytriplet emission with no associated P Cygni absorption featureseveral other emission lines (Stathakis & Sadler 1991; Stritzingerunpolarised emission lines in the outer
2009ip occupies a compact portion of the Q–U plane, in quadrant II.
Meanwhile, a slight increase in P is seen at 5600–5800 Å,
must be emitted by the same region as that with the strong Hαet al. 2012), is not observed or is signiﬁcantly CSM.
weakerin KISS15s
The points form a slightly elongated cluster that shows a continuous
suggesting a potential association with broad He I λ5876 or NaD
emission line (i.e., the CDS and/or photoionized ejecta; Chugai
compared to SN1988Z (Figure 8).
P-Cygni lines.
trend in wavelength, with the bluest points exhibiting the lowest poFig. 4 (top panels) shows the Stokes spectra and the data distribusition angle in the Q–U plane. The cluster as a whole exhibits a small& Danziger 1994; Dessart et al. 2016).
tion in the Q–U plane for the Oct. 14 Lick epoch, binned to 100 Å.
but non-negligible spread in position angle, which corresponds to
The origin of the broad bump at the wavelengths of
4. Discussion
For reference, we also show scaled logarithmic flux spectra in thea range of ∼10◦ on the sky. The dilution in polarization associatedwith intermediate-width Hα emission is apparent as a protrusion
same figure, which has been normalized with a low-order spline
λrest∼7300 Å in KISS15s spectra is unclear, but it may be due
4.1. Host Galaxy Properties
from the cluster of data points which roughly points towards the(avoiding the strongest emission features for the spline nodes). On
to a blend of broad lines of He I and Ca II and several forbidden
origin. Note that our binning causes the line to be unresolved, andOct. 14, near peak brightness, the entire optical spectrum of SN

The Astrophysical Journal, 872:135 (34pp), 2019 February 20

Downloaded from https://academic.oup.com/mnras/article-abstract/442/2/1166/980021 by National Astronomical Observatory of Japan user on 22 October 2018

• 偏光, 非対称な輝線プロファイル (e.g., Hα, Hβ)Outline• Motivation• Method• Simulations and Results• SummarySN 1987A

1984

1987

Gas-radiation coupling in SN explosion• 膨張する超新星エジェクタ• optically thickな状態からthinな状態への遷移• 超新星エジェクタ-星周物質相互作用では, 輻射圧が無視できない

photons

SN ejecta

photosphereGas-radiation coupling in SN explosion• 膨張する超新星エジェクタ• optically thickな状態からthinな状態への遷移• 超新星エジェクタ-星周物質相互作用では, 輻射圧が無視できない
1D semi-analytic model → P01 津名ポスター

CSM

photons

photospheremperature Tr is deﬁned as the temperature obtained from thellowing relation,¶Fri= -G i.
¶t

輻射流体力学シミュレーション
4(6 )E r = a r Tr ,The numerical integration of the basic equt = t n to the next step t n + 1 = t n + Dt ishere ar•is流体方程式、輻射輸送方程式をグリッド状で数値的に解く
the radiation constant. These two values should takecalled operator splitting technique, i.e., inte same value, Tr = Tg, when the energy exchange between
輻射輸送: グレイ近似, 周波数積分されたEr , Frの時間変化
tion and interaction parts one after the oths and •radiation is balanced.
The equations for hydrodynamics can be
We denotethe frequency-averaged
absorption
and scattering
Adaptivemesh
reﬁnement:
超新星エジェクタ全体をカバーしながら細かい構•
some standard techniques for numerical hefﬁcients in the comoving frame by k 0 and s0 . Under theuse the 2nd-order reconstruction scheme
xed-frame
formalism (see, e.g., Mihalas & Mihalas 1984),
造も追う
Leer (1977) to interpolate physical variable
e equations describing the temporal evolutions of hydrocell and the HLLCSuzuki,
schemefor special r
namicalvariables,
the radiationenergythe
自由-自由吸収/放射と radiation
電子散乱• simplestrophysicalJournal,opacity:
00:000000(17pp),
2016
Monthdensity,Day
Maeda, & Shigeyama
namics (Mignone & Bodo 2005) to calc
ux are written as follows,(special
relativistic-)
equency,where
ﬂuxes at the interface of two neighboj¶ (rG00:000000)¶ ((17pp),
rGb )2016 Month Day
e Astrophysical Journal,
Suzuki, Maeda, & Shigeyamhydrodynamics
we
describe
40 )= Grk following,
2 b 2 ) Eour method to nu
+
=
0,
(
7
G
a
T
G
r
(
k
s
G
0 r g00
0r
E r (t , x) = ò In dnd W , j(1 )theadvection
partforradiationand the int
sourceterms¶t¶x
frequency,
where
j
jk
2
3

d

2 b ib j + P d ij )
i)
¶
(
h
G
r
¶ (rhG2bgi , (1 )E r (t , +x) = In dnd W ,
=
G
i
¶x j
¶ti

Fr (t , x) =ò

ò In l dnd W,(2 )2 - P)
2b i)
h
¶
r
G
(
h
¶
r
G
(
g
0, or
the superscript i i speciﬁes thecomponentof =a vectorG+iFr ¶(tsecond
In l dmoment
nd W. Furthermore, theordert, x) = ò¶,x j of the intensity (2)the radiation pressure tensor, iere the superscript ¶iE
speciﬁesof a vector or¶Ftherr component0+ i jorder= - G , of the intensity
sor. Furthermore,
Prij (t , xthe
) ¶=tsecond
In l l¶dxni d Wmoment.
(3 )lds the radiation pressure tensor,
i of radiationij
e dynamical evolution
and gasshould be
¶
F
¶
P
r
rj
i
ij
i
+ In l lj dn=d W-. Gvariables,
Pr (tThe, x) =hydrodynamicald simultaneously.
char- (3)¶tzing the gas are the velocityb¶i ,xthe density ρ, and the gas
ThePgdynamicalgas should, which areevolutionfunctions ofof radiation
the spatialandcoordinates
and beure
transferved simultaneously. Theradiativehydrodynamicalvariables char-òò+ Gr [k 0 - s0 (2G - 1)] b j Fr + G rs0 b j bk Pr ,0 = Grk a T 4 - Gr (k - s G 2 b 2 ) EG(12)0 r g00 2.3.0 Advectionr
for Radiati(8 )2 - 1)] b F j + G 3rs b b P jk ,+ GrThe[k 0 s
(
2
G
j r
k r
0
0 j radiatio
advection equations
for the
i
(12
and
the
are
integrated
by
using
the
s
F
4 ﬂux
i
i
3
i
r
G
=
G
rk
a
T
b
+
G
rs
E
b
0
r
0
r
(9 )
g0
scheme,which
was ﬁrst introduced
by Le
and
i
j
3
i
- order
Gr (k 0 to
+ sintegrate
rs0 badvection
b j Fr0) Fr - 2Gthe
equations
4 ii
3iij rs0 E
3rs-G +=needsTgexpressGr+
r b0 b ib j bkpressure
theGradiation
GGrrk(k00atos00b) b+
P jk .
(tenso
13)
jP +

and

(10)

r

i
jthe followin
3 scheme,itheF-ri . GInr (ks0) Fclosure
0 +M1
r - 2G rs0 b b j Fr

ij
+ Gr2.2.
(k 0 Operator+ sij0) b j P⎛rSplitting. 1 (13
1+-G 3crs0 b ib j b3k P
cjk-

M1 closure Prelationni n=⎜d ij +(The11+)numericalrcode integrates ⎝these2equations by a2 similar
ideal
gas2.2.
way to+thosepresented
inEoS
Takahashi
Ohsuga (2013). At ﬁrst,
Operator&Splitting
we separate the basic equations into the advection andOutline• Motivation• Method• Simulations and Results• SummarySN 1987A

1984

1987

locity is51vmax = 1.3 ⇥ 10 km s .
Fq (#disk ) =Esn = 10 erg. Thus, the corresponding maximum ve 0✓
Zis⇡/2
We consideracentrally
concentrated
CSM
ina
steady
where
F
(#
)
the
coveringqdisk
41
locity is vmax = 1.3 ⇥ 102.2.2.
km sCSM.
) = covering
exp fracti
wind. We assume that the CSM and wind densities fol-Fq (#diskThe ✓Z0 ⇡/2
We consider
centrallya steady
lowatheinversesquareThus,infor
spherical CSMs,
0.310, for #disk = 10
2.2.2.concentrated
CSM law. CSM
F
(#
)=
exp
q
disk
wind. We assume
that profile
the CSMand wind bydensities folthe density
is describedThecovering
fractionyields02Dcylindricalcoordinates(r,z)
2-3dayswith~1000cores•Wea centrally
concentrated
in a steady
2.3.20Ad.
low consider
the inverse
square law.
Thus, forCSM
spherical
CSMs,
0.310, for #disk = 10 andwind.
We assumethatthe⇢(r,
CSM
and
wind
folz) =
⇢wind(R)densities+ 51⇢sph (R),(23) covering fraction yields
Thethedensity
profileisdescribedby◉ SN ejecta with the kinetic energy of 10 [erg]: broken power-law The numerical cod• 10M
low the inverse square law. Thus, for spherical CSMs,
2.3.=forAdaptivemes.
0.310, for #nique
10betterand capt
20
disk
where
the
CSM
are
given by⇢(r, isz)= wind⇢windand(R)by+⇢sphcomponents
(R),
(23)
the density
described
density
proﬁleprofile
The numericalcode is equiptures developingin
2.3.
Adaptivemes
2mericaldomain disco
is co
⇢
(R)
=
A
R
(24) for betternique
capturingwind?given, by(23)
where
the
wind
and
CSM
components
are
⇢(r,
z)
=
⇢
(R)
+
⇢
(R),
CSM
+
wind:
wind
sph
•
The developing
numerical
code
equip
at the lowest
l
tures
in
theisAMR
ejectaand ⇢wind (R) = A? R 2 , with A★=5x10
11 [g/cm]
nique
fordomain
bettercapturing
disco
cessively
coveredbymericalis coveredby
5component:
where
the wind and CSM components are given by (24)• wind ✓◆ptures
in level.
theTheejectaresolutions.
domat thedeveloping
lowestAMR
Then
R
2 2 exp
15 [cm], p=10
⇢
(R)
=
A
R
, withmerical
(25)
and
CSM
component:
Rcsm=5x10
sph
sph
referred
to numerica
as AMR
domainis bycovered
by 5
⇢
(R)
=
A
R
,
(24)
•
cessively
covered
wind
?
R
csm
 ✓
◆p
level
by
1level.
corresponatthe lowest
AMR
Then
resolutions.
The
domain
with
R
2
Mcsm
=0.1M
, 1.0M◉, 10M◉• CSM
andmass:⇢sph(R)= pAsphR
exp(25)Freeparamexpansion:asv=r/t
with
=◉10.
In
the
following,the, wind density
of
2.AMR
The
highestAMreferred tocovered
0 andcessively
bylevelnumerica
R
csm
13
 5✓
eter is set to A? =⇥ 1011 ◆g pcm 1 , whichcorresponds
which
2
times
fine
level
by
1
corresponds
toafin
resolutions.
The
domainwithρ
simulation
startsat
t=1000s
after
the
explosion
R•
2
6
aAsteady
mass-loss
a rateofM ρ∝yrv-m1 (m=1)
for
a Thetohighest
exp theatwind, 10 param(25)level
is AMR
realized.
with⇢sph
p =(R)
10.to=Inthefollowing,
density
sph R
of 2.
is
referredas
AMR
levellevel
0 and
R
1csm
constant
velocity
kmcorrespondss 1 . This value leads
Sincefiner
the resolutioSNeter is set toA? = 5wind⇥ 1011g cm of, 100
which
whichby213
times
level
1 corresponds
to ejecta fin
1 so that it does
athesufficiently
dilutewindcomponent
numberAMR
of numerica
to ap steadymass-loss
at a rate
10 6density
M yr paramfor alevel
is realized.
with= 10. toIn
following,
theofwindof
2. Thehighestlevel is
1
not
a↵ect
SN
ejecta.
The exponenThus,
we decrease
constantwind
100
km 1s, of
.the
This
value
leads
Since213
thetimes
SN ejecta
are
expat
eter
is set toA? velocity=
5 ⇥ the10of11expansion
g cm
which
corresponds
which
finer
resolutio-nρ∝vejecta(n=10)
tial
factorin
thisspherical
density
profilerealizes
6CSM
1it doesastheexpandtoasufficiently
dilutewindcomponentsothat
numberof numerical cellsareto a steady mass-loss at a rate of 10 M yr for alevel
is realized.
a velocitysmoothcuto↵
around, beyond
the we decrease
Alth
1R = Rcsm
not a↵ecttheexpansionof theexponentheare
highestconstant
wind
of
100
kmSN
s ejecta.
. ThisThevalueleads whichThus,
Since thethe
SNsimulation.
ejectaexpa
windcomponent
dominates.
The
CSM
mass Mcsm
istime asa result
factor in
thisspherical
CSM
density
profilerealizesasthe ejecta
expand
and
itare
rea
to tiala sufficiently
dilutewindcomponent
so
that
it
does
numberof withnumerical
cells
spherical CSM
expressed
as aR function
the normalization
imumAlthough
resolved lengtsmootharoundRcsm,ofbeyondthe constantthe simulation.
the snota a↵ect
thecuto↵expansionof the=SN
ejecta.
The whichexponenThus,
wedecrease
the highes
inner
outer
Asph andthe outer The
radiusRcsmmass,
theexpandingejectar orv
wind
component
dominates.
CSM
Mcsm is
with
time
as
aexpandresult
ofand
this
coa
tial factor in this sphericalCSM
density
profilerealizesastheejecta
it
re
15
R
cm
csm=5x10
lowtheevolutionofZ
expressedasafunctionofthenormalization
constant
imum
resolved
lengthcomparea smooth cuto↵ around R = Rcsm , beyond whichthethev=w
simulation.
Although the s
cvej (wc=0.1)
2
t
'
200
days.
M R = 4⇡
⇢ (R)R dR
Aand the outer radius,

SN ejecta colliding with CSM: numerical setups

Spherical CSM: results
spherical CSM with Mcsm=10M◉• 密度,ガスエネルギー密度,放射エネルギー密度の空間分布• 星周物質との衝突による
順行,逆行衝撃波形成• エジェクタ-星周物質間での流体不安定性の発達~5days~10days~20days~50days~100daysInteracting Superno

Spherical CSM: results• 密度,ガスエネルギー密度,放射エネルギー密度の空間分布• 星周物質との衝突による
順行,逆行衝撃波形成• エジェクタ-星周物質間での流体不安定性の発達

SN ejecta
CSM

RT instability

hav
Un
eje
tio
of
the
twe
sev
hyd
the
the
to
add
lon
O
6,
sig
bu
tin
me
she
gas
ach
Th
and
sha
abl

Spherical CSM: results• bolometric light curves.
• CSM小: 早く進化する光度曲線• 球対称モデルでは見込み角依存性は強くない• type IIn SNe: 1998S, 2010jl,
2006gyとの比較

symmetryΘobs
axis

spherical CSM

異なる見込み角からの光度曲線

of Rcsm a
0  raround 6.4⇥10
and, beyond
6.4⇥1016whichcm zthe 6.4⇥1016
mooth cuto↵
R 16=cm
Rcsm
a smooth cuto↵ around R = Rcsm , beyond which the
cm. A
reflection boundary
conditionis csmimposednd component
dominates.
The CSM
mass Mis at the M= 4⇡Adisk (1 + 1
wind axis
component
dominates.
The CSM
mass Mcsm
csmis
円盤状星周物質との衝突
symmetry
r
=
0,
while
free
boundary
conditions
csm
with ✏ = 0.05. M
Thepressed as a function of the normalization constant
expressed
as aotherfunctionof the normalization constantare
imposedfortheboundaries.
is(#expected
to be
aro
ph and the outer radius Rcsm ,
where
F
)
is
the
coveriqdiskAsph Zand the outer(10M◉radius+ 1051 Rerg)
csm ,• 超新星エジェクタは同じ
whereofFqth(ofthe
expansion1
2.2.1. 2SN Z
ejectarespect to the solid angle,1
respecttoenergyrapidly
decrea
M
=
4⇡
⇢
(R)R
dR
csm
sph
ただし、円盤状星周物質を考える
2• 1. Spatial
Figuredistributions
offreelythedensity,thegas anddR
radiationenergy densities, and the r
M
=
4⇡
⇢
(R)R
We initiallyassumeexpandingspherical
supercsm
sph

Z ⇡/2
0
sourceand
immediateshortly aftert =ejecta.
t0 for model
D10_M10.
0
nova
We
start
our
simulationt0 =◉,1000
models
D10:
θ
=10
[deg]
and
M
10M◉s. Fq (#the
disk
csm
•
= 4⇡Asph Rcsm (1 + 1/p),=0.1M◉at, 1.0M
(27)kinetic
counterpa
)
=
exp
disk
=
4⇡A
R
(1
+
1/p),
(27)
sph atcsm
Fq (#dThus, the radial velocity of a layera radius R is ini0 a power-law
velop
with
modelsD20:θdisk=20 [deg]and (1+1/p)
Mcsm=0.1M, 1.0M◉, 10M◉•(x)
ereThus,isfor
a Gammafunctionand'◉0.951
tiallygiven
byspherical
CSMs,
thedensity
profileis dean angle meas
as#anis asymptotic
beha
where (x) is a Gamma function and(1+1/p)where' 0.951
R
p = 10.
The covering
fraction toyield
CSMによるカバー率:
D10
→
~15
%,
D20
→
~30
%
scribedbycomes
negligible
the
•
v= ,
(17)
for p = 10.
The covet0
For a CSM disk, we assume the following
functional
0.310, for#diskTherefore,= 10 #and1989).
from
=
c
For
a
CSM
disk,
we
assumethefollowing
functional
0.310,for= ⇢wind(R),m,
the spatial
distributionsoffor ⇢(r,v <z)vmaxand (R)v =+0⇢sphotherwise.
Here (24)
we denotenal
energy
of
the
eject
form,
the spatia
2
2 1/2
⇢(r,
z)
=
⇢
(R)
+
⇢
(r,
z),
(28)
the
three-dimensional
radius
by
R
⌘
(r
+
z
)
,
in
tion
energydensities,
and
th
wind
disk
that
the
radialdensit
where
the
wind
and
CSM
components
are
given
by
diation energy densities, and ⇢(r,
the radial
velocity(fromleft(r,toz),right) Throught
this tion
work,energ
we
z)
=
⇢
(R)
+
⇢
(28)
wind
disk
order to distinguish it from 2the radial coordinate r. We
ter thet = density
trealized
D10_M10
well
after
thea
0 for modelthisenhanced11ρ (25)
with AA?◆R ,✓1 [g/cm]★=5x10
ter t = t0
◆
 (R)
✓=
also introduce
theinclination
angle
✓
=
cos
(r/R) parameters).
towith⇢windpqtified when the initial
angle ofparameter#disk . T
Adisk an angle measuredR#✓◆p axis✓ z = opening◆qthespecifyfromsymmetry
less likely contr
⇢and
exp
A
R,
# slopeejecta
disk (r, z) =
disk

✓
◆
2
p
of the upper andR#⇢initial
z)
=
exp
,
where
is an
angle
measuredfromtheequator,
csm
disk
0. #R
The
density
structureisassumed
to
be
the
disk
A(r,
R
ity as2.3.
in interactingS
sph2
Adaptive m
R
R
#
csm
disk
⇢commonly
(R)
=
exp
,
(26)
sph
(29)
Althoughweradiative
adopt aflux
ver2
adoptedbrokenpower-lawdistribution
with
⇣
⌘
and
the
R
Rcsm
(29)
r
1
there
would
be roomfor15
the break velocity# v=(Chevalier &. Sokerwith
1989;
MatznerRcsm=5x10[cm], is
p=10,
q=4
(30)
br cosnoradiation
field
inwith p = 10. In the following, the winddensityparamR
ies. For
example,byassu& McKee 1999), 11
simulation,thetherma1eter is set to A? = 5 ⇥ 10 g cm , which correspondsof thewhere
uppertheanddiation,ga
6 q = 4.1 Therefore, height
Throught
this
work,
we
assumeto a steady mass-loss at a rate off310
for a
Mej M yrby thedisk tem⇢ej (r, z) =g(R/t(18) termined
identical,
is immediate0 ), with a half1 the
3
3
thedensityisenhancedaroundregionconstant wind velocity of 100 km4⇡v
s br.t0This value leads
In of
this
study,
we wfistructureCSM
disks
opening
angle
of
#
.
The
parameterqdeterminesthedisk component so that it does
to a sufficiently dilute windenergy ofare produced
andthethfor R < v
t , where M is the total mass of the ejecta they kinetic

円盤状星周物質との衝突• 超新星エジェクタは同じ (10M◉ + 1051 erg)
• ただし、円盤状星周物質を考える• models D10: θdisk=10 [deg] and Mcsm=0.1M◉, 1.0M◉, 10M◉
• models D20: θdisk=20 [deg] and Mcsm=0.1M◉, 1.0M◉, 10M◉
• CSMによるカバー率: D10 → ~15 %, D20 → ~30 %ρ円盤状星周物質との衝突
disk CSM with Mcsm=1.0M◉, θdisk=20 deg• spatial distributions ofdensity, gas energy density,and radiation energy density• SN ejecta are “squeezed” by
the disk CSM.

• strong viewing angle
dependence is expected.

~5days~10days~20days~50days~100days円盤状星周物質との衝突• Lbol=4πRobs2Fout ( Fout: outgoing ﬂux at r=Robs=5x1016cm)• 球対称に比べると若干暗くなる

Interacting Supernovae in 2D• 強い見込み角依存性

Mcsm=1.0M◉

Erad
[erg]

symmetryΘobs
axis

CSM disk

SN ejecta

Lbol[erg/s]rapid rise+slow decay

spherical

slowly evolving

円盤状星周物質との衝突=5x1016cm) in 2D• Lbol=4πRobs2Fout ( Fout: outgoing ﬂux at r=Robs
Interacting
Supernovae

17• 球対称に比べると若干暗くなる• 強い見込み角依存性

Mcsm=0.1M◉

M

=1.0M

Interacting
Supernovae in
csm◉2D

Mcsm=10M◉

17

θdisk=10 deg

symmetry
Θobs
axis

Figure 11. Isotropic equivalent bolometric light curves of the disk CSM models with ✓disk = 10 and Rout = 5 ⇥ 1015 cm. Fromleft to right, we plot the models with the CSM mass of 0.1, 1.0, and 10M (models D10_M01, D10_M1, and D10_M10). The light
curves with viewing angles of 0, 30, 45, 60, and 90 degrees are plotted and compared with 2D spherical CSM models in Figure
8 in the lower panel (thick grey line). In the upper panel, the cumulative radiated energy is also plotted for each viewing angle.

θdisk=20Figure
deg11. Isotropic equivalent bolometric light curves of the disk CSM models with ✓and Rout = 5 ⇥ 1015 cm. Fromleft to right, we plot the models with the CSM mass of 0.1, 1.0, and 10M (models D10_M01, D10_M1, and D10_M10). The light
curves with viewing angles of 0, 30, 45, 60, and 90 degrees are plotted and compared with 2D spherical CSM models in Figure
8 in the lower panel (thick grey line). In the upper panel, the cumulative radiated energy is also plotted for each viewing angle.

CSM disk

SN ejecta

disk = 10

円盤状星周物質との衝突
symmetry axisLsmall viewing angle- steep rise- modest decaytime

CSM disk

CSM disk
CSM disk

SN ejecta

large viewing angle- slow rise
- slow decayLtime

CSM disk

SN ejecta

Slowly-rising type IIn supernovae?
306

A. A. Miller et al.

rise time~400days
type IIn SN 2008iy
Miller+(2010)

4

T. J. M ORI

apparent AB magnitude

Figure 1. Left: apparent optical light curve of SN 2008iy showing the long rise time of the SN, including data from DS (this work), the Nickel tele
22.6work), Catelan et al. (2009) and the CRTS web site. Note, conservatively, we take the start time of the SN (t = 0) to coincide with the first epoch w
5 to 2007 September 13, between the last DS non-detection and the first CRTS
detects the source; however, the true explosion date is likely priorNote that the light curve is very broad, and almost symmetric4about the peak. For comparison, we show the light curve of the standard Type II-P S
22.8
3 which shows many similarities to SN 2008iy at late times (see Section 3.2; SN 1988Z(data from Leonard et al. 2002) and theType IIn SN 1988Z,
2 redshift of SN 2008iy. Right: three panels with images of the field of SN 2008iy. Each image
Turatto et al. 1993), as they would have appeared at the
4 arcmin2 , with north up and east to the left. Black cross-hairs mark the location of the SN. The bottom image, marked (c), shows the last non-dete
the DS data. The middle (b) and top (a)images 1
show the CRTS detections on days 0 and 153, respectively; SN 2008iy can clearly be seen in each of
23.0
SN was not flagged as a transient in either the middle or top image by the CRTS software.

6

Table 1. DS Observations of SN 2008iy.
2009; see Schlegel 1996 for a definition
23.2 of the SN IIn subclass, andFilippenko 1997 for a review of the spectral properties of SNe). The
Date
Mag a
σ mag b
SN was later given the IAU designation SN 2008iy (Catelan et al.
(MJD)2009). Mahabal et al. (2009) notedthat the transient was present on
23.4
CRTS images dating back to 2007 September 13; however, it went
54618.35
17.37
0.12
undetected by the CRTS automated transient detection software
g 17.38
54621.35
0.12
until 2008 because it was blended
r 17.3554623.33
0.12
23.6with a non-saturated, nearby
i 17.34(∼11-arcsec separation) star.
54626.32
0.12
z 17.37
Here, we present our observations and analysis of SN 2008iy,
54627.36
0.12
y54628.29
17.37
0.12
which peaked around 2008 October
23.829 (Catelan et al. 2009) and
54629.30
17.36
had a rise time of ∼400 d. This implies 0that SN
longer200 250 300 350 400 450 500 550 0.12502008iy
100took150
54638.21
17.34
0.12
to reach peak optical brightness than any other known SN. Type II
MJD
57700
54651.20
17.28
0.12
SN rise times are typically !1 week [e.g. SNe 2004et and 2006bp;
a Observed value; not corrected for GalacLi et al. 2005; Quimby et al. 2007a; see also Patat et al. 1993; Liet al., in preparation(a)], and have never previously been observed
tic extinction.
b The calibration uncertainty dominates
to rise "100 d, let aloneFigure400, making
SN
2008iyanotherrare
2. Observed LCs of HSC16aayt. The
dashed lines show theover the statistical uncertainty with values
example of the possible outcomes for the end of the stellar life
time rise
we time,
performed
observations.
The mag,
numbersinof ∼0.12 and ∼0.01
respectively.
cycle. In addition to an extreme
SN 2008iythe
is ofspectroscopic
greatinterest because the uniquethe
circumstellar
medium (CSM) in
figure correspondtowhichthe epochs shown in Fig. 10.
it exploded may provide a link to very long-lived SNe, such as SN
1988Z, and thus provide clues into the nature of their progenitors.

rise time>100days
HSC16aayt
Moriya+(2018)

球対称星周物質でも可能かもしれないが、円盤状だと総質量が少なくて済む→偏光観測などの必要性

Interacting SNe with bumpy light curves?

tric light curves of the disk CSM models with ✓disk = 10 and Rout = 5 ⇥ 1015 cm. Fromhe CSM mass of 0.1, 1.0, and 10M (models D10_M01, D10_M1, and D10_M10). The light60, and 90 degrees are plotted and compared with spherical CSM models in Figure 7 ine cumulative radiated energy is also plotted for each viewing angle.
interacting SNewith bumpy light curves?

symmetry
axis

Θobs

CSM disk

SN ejecta

ure 11. Same as Figure 10, but for model D20_M1.

Interacting SNe with bumpy light curves?

tric light curves of the disk CSM models with ✓disk = 10 and Rout = 5 ⇥ 1015 cm. Fromhe CSM mass of 0.1, 1.0, and 10M (models D10_M01, D10_M1, and D10_M10). The light
iPTF13z(Nyholm+2017)60, and 90 degrees are plotted and compared with spherical
CSM models
in Figure 7 in
iPTF14hls (see, Arcavi+17, Sollerman+18)
e cumulative
radiated energy is also plotted for each viewing angle.
A. Nyholm et al.: The bumpy light curve of supernova iPTF13z

B+2.0
g+0.5
r (P60)−18
R (P48)−Si−1.5

Sollerman et al.: Late observations of supernova iPTF14hls

18

Arcavi+1716This work

interacting SNewith bumpy light curves?

-20
19

17

20

−16

21

−15

19

with bumpy light curves?

20

-18

-17

-16

symmetry
Θobs
21
axis

56

22

Absolute r/R-band magnitudeinteracting SNe

Mr [mag]m [mag]-19

18

Apparent magnitude−17

C o → 56 F e

−14

−13

23

⊙

24

0

100

200

300

⊙

400

500
600
JD − 2456324.98

700

800

900

1000

−12

1100

Fig. 1. Light curves of supernova iPTF13z in the R band from the P48 telescope and in SDSS g, r, i, and Johnson B bands from the P60 telescope.
The B, g, and i magnitudes have been o↵set for clarity, as indicated in the legend. The apparent magnitudes are shown without correction for
extinction. Times for conjunction of iPTF13z with the Sun are marked with symbols. The rate of the 56 Co!56 Fe decay has been drawn forcomparison. The black ticks at the bottom indicate epochs of spectroscopy. The S stands for the 0.26 mag correction of the P48 photometry onto
the-15SDSS system. The absolute magnitude Mr is given for the r band and takes MW extinction into account.

mosaic camera and a Mould R-band filter. The 1.52 m telescope(known as P60) was used with a 2048 ⇥ 2048-pixel CCD camera(Cenko22-14 et al. 2006) and filters in the SDSS (Sloan Digital Sky
Survey) g, r, i (Fukugita et al. 1996), and Johnson B (Bessell
1990) bands.
Photometry of iPTF13z from both telescopes was done using
the pipeline known as FPipe (Fremling et al. 2016). This pipeline
performs
23-13 sky and host subtraction and uses a PSF fit to the SN inorder to perform relative photometry. The resulting photometryis given in the AB magnitude system. A selection of 29 SDSS
stars (Table A.1) in the field surrounding iPTF13z was used as
reference stars for the P48 photometry. No colour terms wereapplied. The host subtraction for P48 R-band images was done
24
-12a template based on 65 co-added images taken with P48
using
CSM disk
S
S
SS
S
SS
between 2009 December 21 and 2010 September 24. This time
range was selected to include the deepest P48 images before de0
200
400
600
800
1000
1200
tection (Sect. 3.2.3). For the P60 photometry, host subtraction in
Rest-frame days since discovery
each filter was done using a mosaic composed of 12 co-added
SDSS frames (Ahn et al. 2012) taken in 2004 April and 2004
SN ejectaof iPTF14hls. The data from the first 600 d, to the left of the vertical dashed line, areJune
Figure 2 Photometric observations
fromasA17.
templates. See Fig. 1 for the iPTF13z light curves. The
We have used the same telescopes for the continued monitoring, as well as NOT, TNG, and HST for the latest epochs. The
fieldverticalof iPTF13z is shown in Fig. 2.
bars on top of the plot indicate epochs of spectroscopic observations, and the “S” symbols in the bottom part indicate theWeepochschecked our P48 R-band photometry against the PTFIDE
of Swift observations. Each light curve was fitted with a combination of two tension splines (dashed coloured lines) to(PTF
provideImagea Di↵erencing & Extraction) pipeline (Masci et al.
continuous representation of the data. On the right-hand ordinate we report the absolute magnitude for the r/R band.2017)
To obtain
photometry, which was computed in the 0–520-day interval.
The mean di↵erence between the PTFIDE output and thethe absolute magnitudes for the BgV i bands, their apparent magnitudes must be shifted by 36.024 mag, 36.019 mag,
36.010output from FPipe was mPTFIDE FPIPE = 0.057 ± 0.069 mag.
mag, and 35.993 mag, respectively.
As PTFIDE is only available for the P48 data, but FPipe is available for both data from P48 and P60, we used the latter for
i-0.8
r/R
V+0.2
V+0.2
g+0.8
B+1.7

⊙

P48
P60-GRBCam
P60-SEDM
LCO-1m
LCO-2mNOT
TNG
HST

ure 11. Same as Figure 10, but for model D20_M1.

iPTF13zN
1 arcminE

Fig. 2. Finding chart for iPTF13z based on the discovery image(Sect. 2.1) from the P48 telescope, taken 2013 February 1.48 UT. Two
black bars point at the SN. The inset image (upper left) is a 1000 ⇥ 1000
colour composite from SDSS DR12 showing the host galaxy with thesame orientation as the main image.

Outline• Motivation• Method• Simulations and Results• SummarySN 1987A

1984

1987

Summary: SN ejecta- CSM disk interaction• massiveな星周物質と衝突する
超新星エジェクタの2D輻射流体

symmetry axisLsmall viewing angle- steep rise- modest decay

力学シミュレーション
time• 円盤状星周物質の場合には強い
見込み角依存性• 観測されるIIn型超新星の多様性を説明できるか？
CSM disk

SN ejecta

large viewing angle- slow rise
- slow decay
LtimeSummary: SN ejecta- CSM disk interaction• 流体プロファイルが利用できる• 星周物質の自由度? (bipolar?

• 偏光の情報shell?)
14• 輝線プロファイル (H alpha, H
beta, etc) とその時間進化.

CSM disk

SN ejecta

Suzuki, Moriya, & Takiwaki