26.12.2011@超新星研究会（京大基研）

クォーク＝ハドロン相転移による超新星爆発での重元素合成

西村 信哉
バーゼル大学（スイス）発表のアウトライン• イントロダクション• クォーク＝ハドロン相転移による超新星とその元素合成環境について• 計算の詳細と結果• 放出物全体の組成• 中性子と陽子過剰の元素合成• まとめ

クォーク＝ハドロン相転移による超新星爆発

重力崩壊・バウンス

QCD相転移ニュートリノ or
磁気回転駆動中心星からエネルギー放出

中性子星（マグネター）クォーク星

クォーク・ハドロン相転移による超新星

neutrino radiative pressure from the proto-neutron
ar surface, i.e. by the so called neutrino driven wind,
own in red. Tracers #020 to #050 have stalled after
e bounce, コアバウンス後に原始中性子星がQCD相転移するシナリオ：
however, they are accelerated and ejected
the neutrino driven wind (displayed in green). Trac4
etejectedal. (2009)、Fischer
et dueal. (2011)
s #051 to Sagert#120 arein a prompt way,toe shock wave
originating from the core bounce(dis- from SN1987A further analysis and improve5
servations
x 10
1.2 blue). We clearly see the division of ments
ayed inmatterof the EoS would be required. The magnitude andthetime delay of the second neutrino burst provide corhich is0.8ejected in a prompt fashion after the 一般相対論的流体、ニュートリノ輸送
bouncerelatedinformation about the critical density, the EoSue), matter
which is coasting and falling in again,
but
0.4
in different phases and EOS
the progenitor
model. bag
For low model
EOS：Shen+MIT
ts reaccelerated
outwardbyneutrino
energydeposition
and intermediate mass progenitor models, the energy of0reen), matter which falls back onto the neutronstar, shock becomes sufficient to drive an explosion
thesecond−0.4
t becomes
part of the neutrino wind ejecta (red),even andin spherical symmetry. The explosion is powered
青：νe into赤：νe gravitational
緑：νμ/τ
by thewill
accretion of matterthe deeper
ally −0.8
matter which stays on the neutron star and123
potential of the more compact PNS, if quark matter isver be ejected
10 (black).
10
10
erg/s]

erg/s]

Velocity [km/s]••

Luminosity [1053Luminosity [10rms Energy [MeV]

radius [km]53

present in the core. The ejecta contain neutron-rich ma1
terial that expands
on a fast timescale and shouldbe
1FIG. 10
3: 5Velocity profiles at different times during the postinvestigated as a possible site for the r-process. With rebounce evolution of a 10 M! progenitor model based on eos1,
spect to the remnant, the narrow range of PNS masses
prompt0illustrating the development of the explosion through different
0.255 0.26 0.265
found in Table I may provide an explanation for the clusTime after bounce [s]stages; thick
solid
line
at
240.5
ms,
dashed
at
255.2
ms,
dash4
teringoftheobserved
neutronstar
masses(gravitational)
10at 255.5 ms and thin solid line at 255.6 ms, dashed at
dottedaround 1.4 M! (see e.g. [21]). The discussed direct black
256.4 ms, dash-dotted at 258.3 ms.
0hole formation at the phase transition could be investigated further 30
in light of the observed connection between
103
supernovae andγ-ray bursts [31]. For other progenitor
25
explosion additionally. We obtain explosion energies ofmodels, a delayed collapse to a black hole may occur after
21050 erg (see Table I). The neutrino luminosities
several10
the explosion20
during the PNS cooling phase.
decrease after the onset of the explosion
(see
Fig.
2
at
∼
ニュートリノ駆動風The presented
15 analysis should be complemented by
350 ms after bounce).
multi-dimensional simulations, to explore the impact of
1
10
In 10
general, the models with eos1 and eos2 evolve in a
known fluid instabilities
thatbe treated0.3
in spher- 0.4
00.1 can not0.2
0.5
qualitatively similar manner. However, the models withical
symmetry.
Anotherinteresting
scenario
would
be
a
0
1
2
3
4
Time after bounce [s]the larger bag constant show a longer PNS accretion timeweakneutrino
drivenexplosion,
followed by a fallbacktime
after
bounce[s]
before the onset of the phase transition due to the larger
Dasgupta
et al.
81 (2010)
induced QCD phasetransition. Since
thePRD
QCD phase
Radius [km]FIG. 1: Neutrino luminosities and rms neutrino energies asmperature.
s reaccelerated outward by neutrino energy deposition
102een), matter which falls back onto the neutron star,
OSYNTHESIS IN THE EJECTAt becomes part of the neutrino wind ejecta (red), and1
10amicmatter
Evolutionof stays
Tracersally
which
on the neutron star and will
0
1
2
3
ver
ejected(black).
latebethenucleosynthesis
evolution of
爆発モデル各層の進化
time after bounce [s]in a postprocessing approach, the dyrequiredin radial Lagrangian mass
0.6
105
on the evolution determined with the
0.5
e AGILE/BOLTZTRAN(see
Mezza4
10993a,b,c; Liebendörfer et al. 2001a)0.4
an adaptive grid, was mapped on a3
130 10mass
zones. This provides the
0.3
lution of physical quantities, such as
ature,10andvelocity of the ejected ma0.2
2tion the neutrino fluxes experienced ニュートリノの影響
0.1
me. The mass zones obtained in this101
as tracer
particles for the remaining
0
0

1

2

3

h are ejected in the explosion
are [s]clastime after bounce
rent categories,
related to their ejec0.6
ermodynamic quantities. As listed in
001 to #100
cover the material from0.5
core to the outer layers with a corre0.4
ge from
1.48000 M! to 1.49482 M! .
参考：8Mn is also displayed in Fig. 1.
0.3
#014
are
ejectedwithin 0.5 [s]Kitauranot
et al.
2006
nce. Therefore
we ignore them in
0.2
processing, plus all matter originat0.1 radii. In Fig. 1, they aret smaller
Tracers #015 to #019 are ejected
Ye

4

Ye

radius [km]爆発モデルからみる放出物と組成

4

0

1

2

3

4

F
(bo
ant
Th

time after bounce [s]Fig.
1.— top: inRadialof mass elements as a funcNucleosynthesis
Electron trajectories
Capture Supernovaetion of time after bounce. The colors indicate the properties ofthese mass elements: black, red, green and blue lines refer to matter which is either (black) not ejected, (red) part of the neutrino
driven wind, (green) initially stalled matter which gets boosted by
the wind and (blue) matter which experiences a prompt ejection.
bottom: Evolution of Ye as a function of time after bounce. The
colors are the same as in the top panel.

F
the
cor
the

As shown in the bottom part of Fig. 1, the blue andred zones are neutronized during the collapse via elec-

超新星シミュレーションと元素合成

r-process （中性子過剰、あるいは高エントロピー）•（通常の）超新星爆発 ( > 10M)Fujimoto 2011
•SASI による爆発•ニュートリノの影響によりYe > 0.48
The Astrophysical Journal, 738:61 (15pp), 2011 September 1

Fujimoto et al.

11

(a)

0.1

9

m/Msunlogρ,logT,logr,Ye

10

1

8

(Fujimoto et al. 2011)•（ONeMg星の）超新星爆発（EC 超新星）
•1Dでも2Dでも爆発)•2D だと Ye > 0.4；weak r-process の可能性
(Wanajo 2009, 2011)•他、磁場や回転などによる超新星•より中性子過剰の物質が放出される可能性あり
7

0.01

0.001

6
5

1e-04

4
11

0.48

0.50

0.52

0.54

Ye (10,000km)(b)

Figure 5. Masses as a function of Ye (10,000 km) of ejecta from the inner region
rej,cc ! 10,000 km for a model with Lνe = 4.5 × 1052 erg s−1 .

9

0.6

8

Wanajo 2011

7

5

4
0.3

0.4

0.5

0.6

0.7

0.8

time(s)Figure 4. Time evolution of density, temperature, radial position, and electron
fraction of an ejecta with (a) Ye (10,000 km) = 0.461 and (b) Ye (10,000 km) =0.559, for a model with Lνe = 4.5 × 1052 erg s−1 . Solid, dotted, dashed, anddash-dotted lines represent the density, temperature, radial position, and electron
fraction of the ejecta, respectively. The electron fraction is shown with a value
multiplied by 10. The density, temperature, and the radial position of the ejecta
are presented in units of g cm−3 , K, and cm, respectively. Ye of ejecta can largely
change only in an inner region near the PNS.
(A color version of this figure is available in the online journal.)

the infall of the ejecta near the cooling region (r < 100 km)
in Figure 4(a). As the density and temperature rise to morethan 109 g cm−3 and 1010 K, respectively, the electron fraction
decreases due to the electron captures. When the ejecta start
to be released via neutrino heating at t = 0.42 s, the electron
fraction increases through the absorption of νe by neutrons.
Figure 1. Snapshotof theconvective
regionof the0.461
2D simulationof an ECSN
Finally,Ye (10,000
km)
becomes
for the ejecta.
at 262 ms after corewithentropynucleon(s; left)onandYe (right).
On bounce
the otherhand,
the per
electroncapturesprotons
are not
Mushroom-shaped
lumpsforof low-Y
ejectedthe early
phaseofe matter are
efficienta proton-richejectawithduring
Ye (10,000km)= 0.559,
the explosion.
as shown in Figure 4(b). This is because the densities of the(A color versioninnerof thisfigure(ris available
in thejournal.)region! 200 km)areonline
relatively
low (!108 g cm−3 )due to the mass ejection during an earlier phase ("texp ). The
electron fraction therefore remains constant and rises from 0.5
expands continuously,
explosion
sets
inkm.
to 0.559 viaandthe νae neutrino-powered
absorption in an inner
region r !
200
at t ∼ 100 ms
in richness
1D andin2D
in thesame
The p.b.
protonthe essentiallyejecta is causedby thesmallwayenergyand with a very
similarenergyνe(∼102008).in our
difference
betweenand ν̄50
For TJankaTν̄eal.
adoptede . erg;
νe and etsimulations, the relation,
4(mn − mpthe) > negative"ν̄e − "νe , holds,
which
In the multi-dimensionalcase, however,entropyleads
Ye > 0.5 (Fröhlich
et al. 2006a),
whereleads
mn andtomap are
profile createdbytoneutrino
heating aroundthe PNSof neutrons and protons, and "ν̄e and "νe are energies
short phase massesof convective
overturn, in which accretion downof anti-electron and electron neutrinos, respectively. We note
flows deleptonize
areforneutrino
heatedthe neuthat "ν̄e strongly,= 15.8 MeV
Tν̄e = 5 MeV
and "nearνe = 12.6 MeV for

Wanajo, Janka, & Müller

0.55

5

1D 2D

0.5 4

0.45

3
2

1e+07

Ye, min = 0.466

6Ye, min = 0.404

Ye(10,000km)

The Astrophysical Journal Letters, 726:L15 (4pp), 2011 January 10

mass [10-3 solar]logρ,logT,logr,Ye

10

0.46

1e+08

1e+09

rmin (cm)Figure 6. Ye (10,000 km) vs. minimum radial position, rmin , of ejecta for a modelwith Lνe = 1
4.5 × 1052 erg s−1 . Ye (10,000 km) is largely different from 0.5 only
for ejecta that fall near the PNS star.

0

Figure 50.40
shows masses 0.45
as a function ofYe (10,000 km)0.55
of
0.50
Ye10,000 km. We find that
ejecta from the inner region rej,cc !
most of the ejecta (98.8%) have electron fractions of 0.49–0.5.
Figuremasses
vs. Yejecta,
1D (blue)
2D (red)explosione for the0.9%
Small2. Ejecta
fractionsof theandand0.3%
in mass,are models.
The
width ofneutron-richa Ye -bin is chosen∆Y= 0.005.
values of Ye
slightly(0.46to<be Y0.49) The
andminimum
proton-riche e<are(0.5
indicatedboth cases.
< Ye for< 0.56),
respectively. Masses of the slightly neutronejecta
areis larger
forinmodels
withlarger Lν ,(Aandcolorproton-richversion of thisfigureavailable
the online
journal.)while the mass fractions of these ejecta are comparable for all
the models.

νp-process：陽子過剰 かつ ニュートリノ放射環境•超新星コア付近の陽子過剰な環境•νp-process (Fröhlich 2006, Pruet 2006, Wanajo 2006)
104

s fraction relative to solar

Crand SecondaryNi
GeEjectaKr
Zr
3.3.CaPrimary
Ti
Fe
Zn
Se
Sr
Mo
1D of2Dthe ejecta with the minimum radial
Electron fraction
position 10
rmin2 ! 200–300 km changes due to high neutrino
flux and/or efficient e± capture (Figure 4(a)). Figure 6 shows
1 as a function of rmin of the ejecta for a model withYe (10,00010
km)Lνe = 4.5 × 1052 erg s−1 . Hereafter, we refer to the ejecta with0 as the primary ejecta, which have high maximum
rmin ! 200
10km
densities "108 g cm−3 and temperatures "1010 K (Figure 3).
-1 heated through the neutrino heating. On the other
The ejecta
10are
hand, the others are refereed as the secondary ejecta, heated103r-process の期待計算結果：r-process 元素合成
元素組成（質量比の対数）

発表資料＠超新星研究会（２０１１年３月）

A（質量数）•“main” r-process は厳しい•“weak” r-process （原子番号〜50への寄与）が起こる

モチベーション•爆発モデルの成功 (Sagert 2009)•ニュートリノでの観測について (Dasgupta 2010)•爆発モデルの詳細解析 (Fischer 2011)•放出元素の宇宙への影響 NEW!
クォーク、ハドロン物理から爆発的天体現象
シミュレーション、元素合成まで
近年、注目されている素・核・宇宙の融合点

例experienced as a function of time.

finally matter which stays on the neutron star andnever be ejected (black).

employ neutrino interactions with matter inTABLE
1clude all weak interactions
affectingthe evozones and their properties
he overallSummary
proton of/ mass
nucleonratio Ye (*give
The neutrino fluxes, resulting from the−2 M ]
zone #
M# [10∆M# [Mcalculation
Ye,NSEtej""]
eutrino-)radiationhydrodynamics
−4
– 014
0.000
− 0.208
1.496
× 10
—
n the001
previous
subsection,
are utilized
to0.20de−5
015 – 019
0.210 − 0.216
1.474 × 10
∼ 0.551.5 ∼e actualrates
asafunction
of
time.
This
is
−5
020 – 050
0.217 − 0.232
1.063 × 10
∼ 0.33∼ 0.5−4om theof1.482
all other
reactionrates
051 –treatment120
0.250 −
1.786 ×10
0.33
∼ 0.50 ∼ 0.5
etermined by local thermodynamic conditionsand M
temperature.
# : mass coordinates relative to the innermost zone of 1.48M"

mass zones: ４つの成分∆M# : the averaged mass of the zone
Ye,NSE : Ye at the end of NSE (below T = 9 GK)
NUCLEOSYNTHESISIN theTHE
EJECTAtej : ejection time after
bounce5

10

4

radius [km]

10

3

10

102

4

Ni

10
.1. Dynamic Evolution of Tracers
They are also shown with respect to their mass as we
The mass zones which are ejected in the explosion are
0
3
4
to calculatethethreenucleosynthesis
evolutionof to theiras Ye -distribution
in1 Fig. 1. 2classified indifferent categories,
relatedtime after bounce [s]ter ejection
within process
a postprocessing
approach,the dy- As listed
and thermodynamic
quantities.
ution
is原始中性子星（放出されない）
requiredradialLagrangianmass
0.6
in Table
1, zonesin#001
to #120,
given withthe ejection
thistimescalereason theevolution
determinedwith Ytheafter
bounce (tmass zoneej ) and the final
e,NSE which
0.9 0.5
0.9
is
Y
at
the
end
of
NSE
(below
T
=
9
GK),
cover
the
ydrocode
AGILE/BOLTZTRAN (see Mezzae
Ye
ニュートリノ駆動風
materialfrom thesurface of theetinnercore at 1.48000
ruenn
1993a,b,c;
Liebendörfer
al. 2001a)layerswith a grid,corresponding
mass ofon1.49482
M! . 0.7 0.4! to
0.7
asedMonanadaptive
was mappeda
原始中性子星からの風

rs which are ejected in the explosion are clasree different categories, related to their ejecs and thermodynamic quantities. As listed in

0.5

0.5
0.2

0.3
0.1

0.3

0.1
0
0

1

2

3

electron fraction, Ye, NSE

Ye-2

prompt（放出質量の大部分〜98%）

0.3

delayed••

NDW

grid of 130 mass zones. This provides thedelayed
成分 quantities, such as
gian evolutionof physicaltemperature, and velocity of the ejected ma第二の衝撃波で放出
in addition
the neutrino fluxes experienced
on of time. The mass zones obtained in thisprompt成分 for the remaining
erred toas tracer particles

NS

•

•

mass, M# [10 M!]

•
•

1

40.1

time after bounce [s]

0
20
40
60
80
Fig. 1.— top: Radial trajectories of mass elements as a
massThezone,# indicate the propert
tion of time after bounce.
colors

爆発モデルと元素合成シミュレーション

爆発モデル•流体力学：AGILE-BOLTZTRAN コード（Liebendörfer et al. 2004）•（球対称）一般相対論的ニュートリノ輸送 ＋ マイクロ物理•EOS：Shen EOS（核物質） + MIT bag model（クォーク物質）•親星：10.8 M by Woosley et al. 2002
元素合成
• 爆発モデルを元にポストプロセス計算• 初期進化は流体シミュレーションの物理量（NSE 状態）• （フル）核反応ネットワーク（4000核種以上）• 反応率一般：REACLIB（Rauscher&Thielemenn 2000）• 質量公式（理論）：FRDM (Möller et al. 1995)• 自発、β-delayed 核分裂• ニュートリノ反応• 放出物質と原始中性子星の進化に応じてダイナミックに計算• 核子との反応のみ考慮（Qian & Woosley 1996）

tron Fermi energies, where electron captures dominate
which
make
matterneutron-rich.
エントロピー と Y ：NSE終了（T = 9 GK）
e-2

mass, M# [10 M!]
0.217

1.482
0.60

90
entropyYe

70
prompt50

delayed

60

0.50
NDW

entropy, sNSE [kB]

80

0.40

40
30

20

40

60
80
mass zone #

100

0.30
120

electron fraction, Ye,NSE

0.209

88

組成分布（α-freezeout）-2

Sr

0.25 0.30 0.35 0.40 0.45 0.50
-2enclosed mass, [10 M!]

ニュートリノ駆動風

mass fraction, log10 X(c) innermost (proton rich)νp-proc.
without0maswith (solid) and withond expansion (boost)
20-50. Bottom: The
drodynamic evolution
reheating around 1.5s, a strong appearancet distribution.

4

0
-1
-2
-3

He
p

56

Ni

57

Ni

52

Fe

48

Cr

-4
0.210 0.212 0.214 0.216-2
enclosed mass, [10 M!]

80

90

Fig. 10.— Overview of dominating conditions, i.e. abundanof a few key nuclei after charged-particle freeze-out, as a fution of radial Lagrangian mass, with variations from (a) typ

組成分布（α-freezeout）

Nucleosynthesis in Supernova Explosions Triggered by a Quark-Hadron Phase Transition
40

60

80

100

120

(a) outer (neutron rich)140

delayedno boost

外側（prompt）

mass fraction, log10 X

20

9

88

0

Sr

90

Zr

4

He

-1

-2

60

Ni

58

Ni

56

Ni

0.50 0.70 0.90 1.10 1.30 1.50
-2enclosed mass, [10 M!]

mass fraction, log10 X(b) inner (neutron rich)

内側
20
40
60
80
100 120 140（delayedmass number, A+ prompt）

op: The comparison of results with (solid) and withthe effect of reheating and second expansion (boost)r the whole range of mass zones 20-50. Bottom: The
wo cases, based on the full hydrodynamic evolution
g (boost) and the neglection of reheating around 1.5
g leads to photodisintegrations, a strong appearance
reshaping of the heavy element distribution.

0

-1

-2

n
4
He
84
Ge
96
Kr
94
Kr
94
Sr
88
Kr
88
Sr
0.25 0.30 0.35 0.40 0.45 0.50
-2enclosed mass, [10 M!]

最終組成：核ゾーン毎に
8

10

内側
（ディレイド）外側（プロンプト）90

130 10mass number, A
50

90

130

10

50

90

130

0
-1
-2
-3
-4
-5
-6

#015

#017

#019

0
-1
-2
-3
-4
-5
-6

0
-1
-2
-3
-4
-5
-6

#020

#040

#045

0
-1
-2
-3
-4
-5
-6

0
-1
-2
-3
-4
-5
-6

#050

#051

#060

0
-1
-2
-3
-4
-5
-6

0
-1
-2
-3
-4
-5
-6

#070

#080

#120

0
-1
-2
-3
-4
-5
-6

10

50

90

130 105090130mass number, A

10

50

90

130

abundance, log10 XA

abundance, log10 XA

ニュートリノ駆動風

50

Nishimura et al.

内側（再加速される成分）

Nucleosynthesis in Supernova Explosions Triggered by a Quark-Hadron Pha

4
Nishimura et al.
再加速（加熱）は光分解により重元素を壊す

80

100

120

delayedno boost

140

3

10

102

10

7

10

103

0

1

2

3

4

0

1

time after bounce [s]

-4

0.6

0

この部分（再加熱）
-1
を無視
9

10

105

101

-3

(a) outer (neutron

13

1011
density [g/cc]

radius [km]

-2

60

104

10

4
2
0

0.3
0.2

0

1

2

3

(b) inner (neutron
0.1

4

0

time after bounce [s]-2Fig. 1.— top: Radial trajectories of mass elements as a func0 time after
20 bounce.
40The colors
60 indicate
80 the properties
100 of120
tion ofthese mass elements: black, red, green and blue lines refer to matter which is either (black) notmass
ejected,number,(red) part of
A the neutrino

140

4time after bounce [s]

1

0.1
0

3

0.50 0.70 0.90 1.10
enclosed mass, [1100raction, log10 X

6temperature [MeV]

0.4

-5

2

-2

0.5

Ye

log10 XA /XA (no boost)abundance, log10 XA-1

40mass fraction, log10 X

0 105 20

0
1

2

3

time after bounce [s]-14with reheating (boost) and the neglection of reheating around 1.5
s.The reheating leads to photodisintegrations, a strong appearanceof 4 He and a reshaping of the heavy element distribution.

最終組成：ニュートリノ駆動風

νp-proc.
without

3

mass fraction, log

overproduction, log10 X/X!

X

典型的な νp-process（ニュートリノ反応なしと比較）

2

1

0

30

40

50
60
70
mass number, A

80

90Fig. 9.— Resulting nucleosynthesis for the entire mass zones in

F
of
tion
exp

Ye (standard)

0.20

0.30

0.40

0.50

Ye - 10%

-1

-2

-2

Ye - 20%

Ye - 30%

-3

-1

-1

-2

-2

-3

-3

0.20

abundance, YZ

0

0.30
0.40
0.50
electron fraction, Ye

0.20

0.30
0.40
0.50
electron fraction, Yeabundance, log10 YA

0.50

abundance, log10 YA

mass fraction

0.40

-1

-3

mass fraction

0.30

mass fraction

0.20

0

electron fraction, Ye

mass fraction

electron fraction, Yeabundance, log10 YA

integrated ﬁnal abundances

abundance, lo

However,capwe expect that any uncertainties beyond 10%
ak interactions (electron captures and neutrino
are unrealistic. This would not change the discussion of) which determine Ye . For this reason wethe also
results regiven above.
d the present nucleosynthesis calculations with vari-5
40
80
120
160
200
s in the initial Ye for the mass zones experiencing-2
standard
pt andexplosions, according to the recipe
Ye delayedの不定（依存）性
pcor = 10
pcor = 30
!
"
-6
pcor
pcor = 40-3Ye, cor = 0.5 + (Ye − 0.5) × 1 +
,
solar
100
-4
e Ye, cor ’s are corrected ones and pcor denotesthe-7ntage of uncertainty in deviations of Ye fromthe-5
metric value 0.5 and enlarges these deviation fromThe nucleosynthesis results are also presentedin-3-6
11 and show the options of obtaining a full r-process.
-4
ever, we expect that any uncertainties beyond-7 10%
nrealistic for the explosion model we adopt in the-5
NDW
nt work. Thought we assume 10% differences -3would
delayedhange the discussion of the results given above,for-6 prompt-4
ationsbeyond 10% of uncertainties we need
※純粋に本モデルに基づいた-5 dif-7
3
t progenitor,
explosion modelsstandard
or additional -6input
Ye - 10%不定性としては１０%程度
2cs i.e.
multidimensional hydrodynamics
and -7differYe - 30%-8
Ye - 40%
EOS’s. 1-8
HD122563-9
is obvious
0 from the discussion above, that only a-9各成分毎で比較→k” r-process
can be supported by the nucleosyn-1
120 40 160200
80s conditions
found in the explosion mechanism40dis- 80-2mass number, A0abundance, log10 YZ

features. One could argue, however, that such observed
abundance features are a combination of at least two polheintegrated
ﬁnal
abundances
lutions,one(lowlevel)
solar r-contribution plus anotherhe”weak” r-contribution extending only up to A = 130.
m
in 観測との比較（太陽系、weak r-process）
s.
%
standard
0
pcor = 10
ofpcor = 30
pcor = 40
HD122563

-1
-2
-3
-4
-5
30

40

50

60

atomic number, Z

70

80

まとめ

• r-process 元素合成について• reproduce A〜100 r-element (weak r-process)• main r-process のためにはは Y が４０％マイナス• ( 不定性の範囲内でも main r-process はムリ )• 別の親星、EOS、多次元モデル etc. があれば。。。
• ニュートリノ駆動風について• 今回の爆発モデルだと、通常の超新星と同様の環境• A 〜 60 を越えて A 〜 90 の陽子過剰安定核を生成• EOS によって Y の進化やニュートリノ物理量ee（光度、平均エネルギー）変わってくればおもしろい