「超新星爆発と数値シミュレーション」@基研 2011.12.26-28

ANRV284-AA44-12

ANRV284-AA44-12

Annu. Rev. Astro. Astrophys. 2006.44:507-556. Downloaded from www.annualreviews.org
by University of Tokyo on 06/17/11. For personal use only.

Aspherical Nature
Hypernova: Eex~1052 erg
SN1998bw→ 非球対称爆発モデル(Maeda&Nomoto 2008)
GRB: SN Ic との併発例

28 July 2006

28 July 2006

14:21

14:21
Figure 2
The location of SN 1998bw
as viewed at late times by
Chandra (top) and HST(bottom). The circles
represent the 1σ
astrometric position from
Chandra. The fading X-ray
source (“S1a”) consistent
with SN 1998bw is to thesoutheast (bottom left). From
Kouveliotou et al. (2004).

! = (1 − β 2 )−1/2 . This does not necessarily require that the SN be bright, or even
exceptionally energetic, though GRB-SNe often are. It also does not preclude theexistence of SNe without GRBs, powered by the same energy source (Section 3.4.2).
But to produce a GRB, one needs at least as much energy in relativistic ejecta as
is observed in γ -ray and afterglow emission. That is, ERel ! Eγ . The value of Eγ isdifficult to measure directly because of the effects of beaming, but in typical bursts, it isaround 1051 erg (Frail et al. 2001; Bloom, Frail & Kulkarni 2003). ERel can be inferredfrom radio observations at such late times that beaming is no longer important, and is∼5 × 1051 erg (Berger et al. 2003c; Berger, Kulkarni & Frail 2004). Of course, there
can be considerable variation in both these numbers.
The SNe accompanying GRBs also differ from common SNe (Filippenko 1997)
in other ways (Table 1), most obviously the absence of hydrogen in their spectra:
GRB-SNe appear to be Type I SNe. Indeed, where spectra of sufficient quality exist
to be sure, the SN is of Type Ic-BL. Some of the broad peaks seen in the GRBSNe spectra are likely due to low opacity, rather than due to emission from a single
ion spread over large velocity ranges (e.g., Iwamoto et al. 2003). Near maximum
light, GRB-SNe do appear to show broad absorption lines of O I, Ca II, and Fe II

Annu. Rev. Astro. Astrophys. 2006.44:507-556. Downloaded from www.annualreviews.org
by University of Tokyo on 06/17/11. For personal use only.

SN Ic• H, He Emission Lineがない• Core-Collapse• CO Wolf-Rayet 星 の爆発?

ARI

ARI

Woosley & Bloom 2006	

Table 1 Properties of good candidate supernovae associated with γ -ray bursts, X-ray flashes, and X-ray rich
γ -ray bursts(Section 2.1.6) are now considered the closest cosmological instance of GRB 980425

Name
Burst/SN

and SN 1998bw (Soderberg et al. 2004b).

z

Peak [mag]SN likeness/
designation

a
Tpeak
[day]

2.1.5. Late-time bumps. Viewing GRB 980425, and its origin, as distinct from

References

Ic-BL

b

contemporaneous SN
would be a naturalconsequencea massive star 10
origin(e.g.,
GRB 030329/2003dh
0.1685MV =−18.8 toof−19.6
– 13

Ic-BL

c

Ibc-BLd

was interpreted M
as being
caused by a coincident SN∼15
at about a
XRF 020903& Kulkarni 1998) 0.25
V = −18.6 ± 0.5

Ic-BL

e

I (IIn?)f

Ig

GRB 980425/1998bw
Mthe−19.16±Though0.05 the detection
17
V =the “cosmological’’0.0085
set of GRBs wasnorm
in 1998.
of aWoosley 1993; Hansen 1999), no other GRB had obvious late-time emission that

GRB 031203/2003lw
MV = “bump’’−19.0 following
to −19.7
18 –(Bloom
25
resembled a SN. A0.1005
report of a red emission
GRB 980326
GRB 011121/2001dk

0.365

GRB 050525a

0.606

GRB 021211/2002lt

1.00

GRB 970228

0.695

XRR 041006

0.716

XRR 040924

0.859

GRB 020405

0.695

a

MV = −18.5 to −19.6

12 – 14

MU = −18.4 to −19.2

∼14

Ic

h

MV = −18.8 to −19.5

MV ≈ −18.8

12

MV ∼ −19.2

∼17

I

i

16 − 20

I

j

∼11

?

k

∼17

I

l

www.annualreviews.org • Supernova–Gamma-Ray Burst Connection

MV = −17.6
MV ∼ −18.7

515

The time of peak brightness is reported in the rest frame if the redshift is known, observed frame otherwise.
Galama et al. 1998a.
c
Hjorth et al. 2003; Stanek et al. 2003; Bloom et al. 2004; Lipkin et al. 2004.
dMalesani et al. 2004; Cobb et al. 2004; Thomsen et al. 2004; Gal-Yam et al. 2004.
e
Soderberg et al. 2005.
f
Bloom et al. 2002; Garnavich et al. 2003b; Greiner et al. 2003b.
b

A&A 512, A70 (2010)#%%*D#%%'4E#%%&.F"$$"6<"$$+G:"$$'.7"$$'.9!"#SN2007bi (Gal-Yum+2009)• 非常に明るいSN Ic• 大量の56Niを生成する (>4Msun)
56Ni→56Co→56Fe• [OI]の輝線より球対称に近い爆発(Young+2010) 	

)(

B0,

-./01234,5678932:4!#%30, )(
C4

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!#*

Young+2010	

!#+

,$

,)$

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,"$$

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,+$$

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,$

,"$

,*$

,($

, , , , , , , , , , , , , , , , , , , , , , , , , , , , , , ;<0=>,?:6@/A

D. R. Young et al.: Two type Ic supernovae in low-metallicity, dwarf galaxies: diversity of explosions   



!!#"



Fig. 5. R-band lightcurve comparisons of SNe Ic (V-band used for SN1997ef comparison). All comparison data have been corrected for extinctionand epochs are relative to the R-band maximum light as stated in the literature and given in the rest frame of the individual SNe. References andcolour-excess: SN 1994I with E(B − V) = 0.300 (Richmond et al. 1996), SN 1997ef with E(B − V) = 0.042 (Iwamoto et al. 1998), 1998bw withE(B − V) = 0.065 (Galama et al. 1998; McKenzie & Schaefer 1999; Sollerman et al. 2000; Patat et al. 2001), SN 2002ap with E(B − V) = 0.09
(Foley et al. 2003; Yoshii et al. 2003; Pandey et al. 2003; Tomita et al. 2006), SN 2003jd with E(B − V) = 0.14 (Valenti et al. 2008a). The panelon the right provides an expanded view of the early-time behaviour of the lightcurves.
!!#"

!'#$

%::!;
%::<=>&""&84&""'?7&""@,9&""A=,&""A=B!'#$[OI] のDoublepeakは非球対称効果	

Young+2010	
Fig. 15. A comparisonof the profiles of the [O ], Mg ] and [Ca ] emission lines found in the nebular spectra of SNe 1998bw (black), 2003jd








20

!&#$

I

I

16

14

12

10

!'#"

II

MORIYA ET AL.

!&#$

!&#"
-21

Table 1
Amount of Elements Contain

CC100 (no18mixing)
CC100 (half mixing)
CC100 (full mixing)
16
Y10 (bolometric)
G09 (with bolometric correction)

Na I λ5891
Si II λ6355
O I λ7774
Ca II near-IR λ8571

!%#$
-20

absolute bolometric magnitude

Ejecta Velocity (1000 km s-1)

18





!%#"
-19
!"#$
-18 ($"

)"

)$"

)%""

)%$"

)&""

)&$"

Ejecta Velocity (1000 km s−1)(%*+,)-)./0,)1 2!'#"(grey) and thefinal +367 d spectrum of SN 2007bi (red).
L84

Core Collapse シナリオ (Moriya+2010)
43Msun のCO starの球対称爆発
56Ni量やLight Curveを説明可能!"#   Pair Instability Super Nova シナリオ
~100Msun He星のPISN○Yieldを再現する
×He外層が残っている

!"$

)'""

12 C

1.4

14

10
)'$"

O

Ne

24 Mg

28 Si

18.7

1.4

1.5

5.1

Note. Units: M!!%#$.

12) ) ) ) ) ) ) ) ) ) ) ) ) ) ) ) ) ) ) ) ) ) ) ) ) ) ) ) ) ) 34+56).78912

Ca II H&K λ3976
Fe II λ5037
Si II λ6355
Ca II near−IR
λ8571
16
20
!&#"
O I λ7774 (limit)

(&"

)"

)&"

)!"

)@"

3. CORE-COLLAPSE SN MOD

8

Fig. 6. BVRI-17
quasi-bolometric lightcurve comparisons. All comparison data have been corrected for extinction and We
epochsconstruct
are relative toseveral
thecore-collapse
R-band maximum light and are give in the rest-frame of the individual SNe. References and colour-excess: as in Fig. 5, also SN 2006gy with68
them
the observations of the bE(B− V) = 0.56 (Agnoletto et al. 2009). The panel on the right provides an expanded view of the early-time behaviourof thewith
lightcurves.

velocities of SN 2007bi shown in Y10
not cover the rising part of the LC, w
magnitude of the rising part from the R
2
4
50
60
70
80
90
100
110
120
130
140
0 6 of 19-15
10
20
30
40
50
60
Page
the same bolometric corre
Epochassuming(days)
-100
100
200
300
400
Epoch0 (days)
R-bandmaximum.
We
also take intodays since the R band maximum magnitude
Time given in the SN restFig. 16. Ejecta velocities of SN 2007bg. Time given in the SN Fig. 17. Ejecta velocities of SN 2007bi.
possiblethatajetemergesfrom the cframe.
Note
that
the
velocity
of
O
I
λ7774
measuredinthespectraofrest-frame. Figure 1. Bolometric LCs of the C+O star SN models CC100 (Mej = 40 M! ,SN 2007bi can only be treated as a lower-limit;
a consequenceof con- (e.g., Maeda &the mixingof the ejecta
Ekin = 3.6 × 1052 erg, and M56 Ni = 6.1 M! ). The tamination
observed inbolometric
LCby the peculiar [Ca II] feature.
the blue-wing2009).
(open circles) is taken from Y10. The bolometric magnitude of the rising part
The LCsof ofof SN 2007bi (open square) is estimated from the R-band
magnitude.
Allusedtheto measure abundances
Table 8.
Emission lines
the successful
host galaxies models (CC6-16Moriya+2010	

4

1

Eex =

Spherical Explosion
爆発エネルギーの解放(Eex)
衝撃波球面の温度と半径(Radiation Dominant)	

4πa 3 4Eex =R T
3
4πa 3 4Eex =R T
356Niの合成温度	
4πa 93 4TENi= 5 × 10
R TK
ex =
3
TNi = !
5 × 109 K
3 9EK
ex
35 × 10
TNiNi==!
R
4
56Niの合成領域	
4πa
T
Ni
3
E
ex
RNi = 3!
4πa TNi 4
3 Eex
RNi = 34πa TNi 4

4πa 3 4R T
3

(1)

TNi = 5 × 109 K
!

(2)
11

3 Eex
4 2, 2008
No.
4πa
4πa T3Ni
4πa
3 44
EEexex== 1 3 RR TT
3
1

RNi = 3

(3)

SYNTHESIZED MASS OF 56Ni

(1)
(1)

=551××10
1099 KK
Ni =
TTNi
(1)
!
!
(1) 33 EEex
ex
= 33
Ni =
RRNi
44
4πa
(2)
Ni
(1)4πa TTNi
(2)

(2)
(2)

(3)
(3)

Umeda&Nomoto2008	

(2)
(3)
(3)

sults, especially th
16
O, are applicablecore masses. Usin
simulated the SN
the mass of the sy
We find that the
Si-burning stages.
of the stars longer
core growth durin
massive stars can
existence of thislarge M(56 Ni).
We have shownwith the increasinmass M. Among t
energy because mcurves in the dens
increases steeply w
larger 56 Ni produc
this also leads to aa large Eexp doesM ( 56 Ni) ¼ 2:2, 2
30, 50, 80, and 10
MCO ¼ 11:4, 19.3
It is possible to p
if MCO # 34 M"
M ( 56 Ni) k 13 M"
collapse SN mode
most massive modplosion energy of
13 M", but such E
is no observationa
Nevertheless, w
results to more ma
as 60 M". Such a m
M ( 56 Ni) # 13 M"
we are constructin
results in a forthco
stars above 300 M
Those stars may beholes (e.g., Fryer e
may not accompan#10 M 56 Ni (OhFig. 7.— Upper limit of 56 Ni masses produced by core-collapse SNe withvarious masses and explosion energies. 56 Niup shown here are the ejected 56 Ni
mass when the mass cut is located at the top of the Fe core. If mass cut is larger theejected mass will be smaller.

How to produce
much 56Ni ??
(3)• Eex を上げる	
• プロジェニターサーチ （星の進化計算）
RNi 内に含まれる質量を増やす
models which do not produce large Zn/Fe (UN02). In the energetic explosions Co/ Fe ratios are also relatively high. However, as shown in UN05, the Co/ Fe ratio is not much enhanced
for Ye P 0:5 even for the energetic explosions. In the present
models, indeed, the Co/ Fe ratios are less than the solar value.
More detailed study about nucleosynthesis and its parameter dependences are given elsewhere.
4. CONCLUSIONS AND DISCUSSION
Motivated by the discovery of unusually bright SNe 2006gy
and 1999as, we have investigated how much M(56 Ni) can be synthesized by core-collapse massive SNe with the initial mass of
M ! 100 M". Observed properties of SN 1999as and SN 2006gy
suggest that the ejected 56 Ni masses are M ( 56 Ni) # 4 M" (Hatano
et al. 2001) and 13 M (Nomoto et al. 2007). Previously the only

A progenitor for SN 2007bi

MCO/M!

Mf/M!

Progenitor Model• M = 110M
Z=0.004• Mass LossによりHe層まで失ったCO Wolf-Rayet星→ M = 43.1M
Yoshida&Umeda 2011	

180
160
140
120
100
80
60
40
20
0

(b)
Yoshida&Umeda2011	
PISN

CCSN

100 150 200 250 300 350 400 450 500
MMS/M!

Figure 1. The final mass (a) and CO core mass (b) as a function of theMS mass. Solid, dashed, and dotted lines correspond to cases A, B and C,
respectively. Squares, triangles, crosses and circles indicate WN, ‘He-rich’

The range of MS mass for CC SN models is shown as the lightshaded region in Fig. 1(b). The range extends from 100 to 280 M!
in case A. In case B, all models except for MMS = 100 M! will
explode as CC SNe appropriate for SN 2007bi. On the other hand,in case C, the mass range is limited to 100 ! MMS ! 170 M! .

Hydrodynamics• Code: 2D Axis-symmetric3.3• Explosion
Surface He abundance
for SNe Ib/Ic
Energy:Eex= 30×1051 erg
SNe
Ic are•0< characterizedt < Tend by=weak
500or absent
sec He spectra. However,a quantitative criterion to distinguish between SNe Ib and Ic has
not been theoretically established. The He lines are considered toappear because of the excitation of He by non-thermal electrons
excited by γ-rays from the decays of 56 Ni and 56 Co (Lucy 1991).
100
The strengthof the He lines should be sensitive to the amounts ofHe and 56 Ni, the amount of matter in intermediate layers between
Ni and He layers which attenuate the γ-rays, the degree of mixing of Ni into He layer, etc. Criteria to distinguish between SNe
Ib and Ic were discussed using the total He others
mass of a progenitor(Wellstein &10
Langer 1999; Georgy et al. 2009; Yoon et al. 2010)or the He mass fraction at the outermost layers (Yoon et al. 2010).
Fall-backThe effects of thickness of
the intermediate layers and the degree
56 1997). We
of mixing have been investigated (Woosley & EastmanNiM
discuss the possibility of
cuta SN Ic progenitor by considering the total He mass, the He mass fraction at the surface and the mass ratio
1of He to the intermediate
layers.
0
5
10
15
20
25
30
The first criterion is based on the total51He mass. We consider[10Ic: erg]
twocases forExplosion	the He mass limit E
ofexSNe
0.5 and 1.5 M! . In
Spherical
previous studies, the He mass limit of SNe Ic was assumed to be
0.5 M! (Wellstein & Langer 1999; Yoon et al. 2010) or 0.6 M!
(Georgy et al. 2009). On the other hand, Georgy et al. (2009) reported that the choice of the total He mass limit between 0.6 and1.5 M! hardly affects the MMS ranges for SNe Ib/Ic. Yoon et al.
(2010) suggested in discussion that He lines are not seen in early-[Msun]

180
160
(a)
140
120
100
80 MS
sun,
60
40
20
0
CO150 200 250 sun
100
300 350 400 450 500
MMS/M!

3

2

Model & Method

We performed spherically symmetric and aspherical explosion of massive CO star employing 2D special
relativistic hydro-dynamical code. The progenitor model is one of investigated in Yoshida+2011 which isthe end product of M = 110M! and Z = 0.004 star in main sequence. It loses the H and He envelope bymass loss wind until the pre-supernova stage and becomes 42.7 M! of bare CO core. Explosion energy
is adjusted to 30 × 1051 erg and injected from just outside from 2M! of mass-cut during 0.1 sec of the
beginning of calculation. Precedent studies (eg. Umeda+2008, Moriya+2010) show that explosion of
Aspherical51∼ 100M! star witharound 30 × 10Explosion
erg induces ∼ 3 − 6M! of 56 Ni in spherical explosion. Now weconcentrate on theAbrupt
asphericalenergy
effects of injection
explosion of massive CO star on produced 56 Ni. Summarized inTable1, calculated models have various opening angle θop . Hydro-dynamical evolution is calculated until
Opening Angle (θop)homologousの異なる7モデルeop
t = 500 sec when the velocity becomes almost
as v ∝ r.
Name
θopθop /90◦

OA1
7.03◦
5/64

Table 1: Calculated models
OA2 OA3 OA4 OA5
11.3◦ 22.5◦ 45.0◦ 67.5◦
1/8
1/4
1/2
3/4

OA6
78.8◦
7/8

OA7
90.0◦
1

inner boundary

After the hydro-dynamical simulation, we also employed postprocess for calculating nucleosynthesis.
Thermodynamical hystories are traced by maker particles. Nθ × Nr = 40 × 130 = 5200 particles initiallycover the whole progenitor,
and typical mass resolution is ∆M ∼ 0.001M
star.
! around the center of the
Nucleosynthesis
Thermodynamical
Histories• Post-process計算
Results • 282核種のNetwork (n,p ~ Br)

Tracer Particle Method
3• 5200 particles:(Δm)
Fig.1 is the final distribution of each chemical species in velocity space for OA3
model.
It showssun
thei~0.001M

significantly aspherical feature that 56 Ni is ejected only polar direction. This is because injected bubble
with high temperature goes upward along the polar axis, and then, 56 Ni is produced around hot bubble
and ejected. To the contrary, stellar matter in equatorial region cannot avoid accreted onto the central
remnant. 28 Si is produced in lower temperature than 56 Ni, so distributed around 56 Ni. As 56 Ni is dissolved
into 4 He via photodissociation under very high temperature, 4 He exists only around the bubble.
Fig.2 shows the ejected mass of the elements, 4 He, 28 Si, 56 Ni and total ejecta for all models. The
global feature is clearly seen that each ejeced mass decreases with smaller θop . This can be understoodEjected Mass [Msun]
5

Ni
28
Si
16
O
Ejecta

40

3

30

2

20

1
0
0°

10

OA5 OA6
45°
Opening Angle

90°0O, Ejecta

56

56

16

Ni, 28Si

450

{X j i } (Σj X j i = 1)

Simple Evaluation of 56Ni	

Ti (∀ tn ) < TNi → X 56Ni i = 0

5

0.5

4

0.4
Fall-back
0.3

2

0.2

56

3

Produced
1
OA5
0.1
Ejected(post-process) Ejected
OA6
Fall-back/Produced
0
0
0°
45°
90°
Opening Angle

Fall-back/Produced 56Ni

Ni Mass [Msun]

Ti (∃ tn ) > TNi → X 56Ni i = 1

Ejected Mass [Msun]
5

Ni
28
Si
16
O
Ejecta

40
30

2

20

1

45°
Opening Angle

90°

0

5

0.5

4

0.4
Accretion
0.3

2

0.2

56

3

1

Produced
Ejected(post-process) Ejected
Accretion/Produced

0
0°

45°
Opening Angle

OA5

0.1
OA6
90°

0

Accretion/Produced 56Ni

Ni Mass [Msun]

10

OA5 OA60
0°O, Ejecta

3

16

Ni, 28Si

4

56

50

56

M(ejec) = 17.4 Msun
M(56Ni) = 1.77 Msun

5e+09
56
Ni
othersaccreted

4e+09

Z [cm]

3e+09

2e+09

12Msun

7Msun

1e+09

2Msun
0

1e+09

2e+09

3e+09
R [cm]

4e+09

0
5e+09

Ejected Mass [Msun]
5

Ni
28
Si
16
O
Ejecta

40
30

2

20

1

45°
Opening Angle

90°

0

5

0.5

4

0.4
Accretion
0.3

2

0.2

56

3

1

Produced
Ejected(post-process) Ejected
Accretion/Produced

0
0°

45°
Opening Angle

OA5

0.1
OA6
90°

0

Accretion/Produced 56Ni

Ni Mass [Msun]

10

OA5 OA60
0°O, Ejecta

3

16

Ni, 28Si

4

56

50

56

2e+09Z [cm]

12Msun2e+09

12Msun

7Msun

7Msun

2Msun

1e+09

2e+09

2Msun

0

3e+09

1e+09

Ni Accretion Rate [Msun/sec]
0.12
0.1
0.08
0.06
0.04
0.02
0

0.3
0.2
0.1
0

20

40

60
80
time [sec]

2e+09

0

3e+09

0

1e+09

R [cm]

56

0.4

0

1e+09

0

R [cm]

56

2e+09

12Msun

7Msun

0
0

3e+09

1e+09

1e+09

2Msun

56
Ni
othersaccreted

3e+09

Z [cm]

56
Niothersaccreted

3e+09

Z [cm]

56
Niothersaccreted

100

120

140

M(ejec) = 28.7 Msun
M(56Ni) = 2.42 Msun

2e+09

3e+09

R [cm]

56

Ni Accretion Rate [Msun/sec]

Ni Accretion Rate [Msun/sec]

0.1
0.08
0.06
0.04
0.02
0

20

40

60
80
time [sec]

100

120

140

M(ejec) = 36.8 Msun
M(56Ni) = 2.47 Msun

0
0

20

40

60
80
time [sec]

100

120

140

M(ejec) = 40.2 Msun
M(56Ni) = 4.01 Msun

Chemical AbundanceEjected by Aspherical Explosion of Massive CO Star
• 中心付近で生成される isotope は θop に依存して Fallback や流体力学的効果を特に強く受ける• 56Ni は θop/90° ≥ 7/8 で 4Msun以上放出→ 観測により示唆されるSN2007biの球対称性は 偶然ではない