Physical ingredients of corecollapse supernova driven by
neutrino-heating mechanism
Yudai SUWA

(YITP, Kyoto University)
Collaboration with

T. Takiwaki, K. Kotake (NAOJ), M. Liebendörfer, S.C. Whitehouse (U. Basel), T. Fischer (GSI),
S. Yamada (Waseda U.), K. Sato (NINS & U. Tokyo)

Core-collapse supernovae
One of the most energetic explosions in the universe
Eexp〜1051 erg
Egrav〜1053 erg (〜0.1 M⦿ c2)
Eν〜1053 erg

Formation of neutron star / slack hole
Formation of gamma-ray bursts?
✤ All known interactions are important•Microphysics•Macrophysics
Gravity

Weak

Elecromagnetic

Strongneutrino physics

core collapse
pulsar, magnetar,magnetorotational explosion

Oct 29th 2013

equation of state of dense matter

Supernovae and Gamma-Ray Bursts 2013

2 /21

Explosion energy
Tanaka+ 09

100

51

Kinetic Energy (10 ergs)

(a)

03dh

02ap
06aj

1

93J
94I

87A

Hypernova Branch

05bf
99br
97D

0.1

10

15

03lw

97ef

08D

10

98bw

Faint SN Branch

20
25
30
35
40
Main Sequence Mass (M )45

Observationally,
there are several
classes of SNe- Superluminous SNe- Hypernovae/GRBs- Ordinary SNe- Faint SNe

50

Where is the upper limit of explosion energy obtainable
by neutrino heating mechanism?
Oct 29th 2013

Supernovae and Gamma-Ray Bursts 2013

3 /21

Systematics in supernova simulations
Our Goal: Produce Successful Explosion! of ~1051 erg
Dimensionality of hydrodynamics

Iwakami+ 08, Nordhaus+ 10, Hanke+ 11,
Takiwaki+ 12, Couch 12, Ott+ 13

General relativity

Liebendörfer+01, Müller+ 12, Kuroda+ 12

Neutrino physics
Scheme to solve Boltzmann equation

Ott+ 08, Shibata+ 11, Sumiyoshi & Yamada 12

Interaction rate

Langanke+ 03, Arcones+ 08, Lentz+ 12

Collective oscillation

Raffelt & Smirnov 07, Duan+ 10,
Dasgupta+ 10

Nuclear equation of state
Initial condition
progenitor structure (mixing, wind...)Lattimer & Swesty 91, H. Shen+ 98, G.
Shen+ 10, Furusawa+ 11, Hempel+ 12
Nomoto & Hashimoto 88, Woosley &Weaver 95, Woosley+ 02, Limongi & Chieffi
06, Woosley & Heger 07, Yoshida+ 12

rotation / magnetic ﬁeld
Oct 29th 2013

Supernovae and Gamma-Ray Bursts 2013

4 /21

Systematics in supernova simulations
Our Goal: Produce Successful Explosion! of ~1051 erg
Dimensionality of hydrodynamics

Iwakami+ 08, Nordhaus+ 10, Hanke+ 11,
Takiwaki+ 12, Couch 12, Ott+ 13

General relativity

Liebendörfer+01, Müller+ 12, Kuroda+ 12

Neutrino physics
Scheme to solve Boltzmann equation

Ott+ 08, Shibata+ 11, Sumiyoshi & Yamada 12

Interaction rate

Langanke+ 03, Arcones+ 08, Lentz+ 12

Collective oscillation

Raffelt & Smirnov 07, Duan+ 10,
Dasgupta+ 10

Nuclear equation of state
Initial condition
progenitor structure (mixing, wind...)Lattimer & Swesty 91, H. Shen+ 98, G.
Shen+ 10, Furusawa+ 11, Hempel+ 12
Nomoto & Hashimoto 88, Woosley &Weaver 95, Woosley+ 02, Limongi & Chieffi
06, Woosley & Heger 07, Yoshida+ 12

rotation / magnetic ﬁeld
Oct 29th 2013

Supernovae and Gamma-Ray Bursts 20134 /21

1D simulations: fail to explodeRammp & Janka 00

Liebendörfer+ 01

By including all available physics to simulations, weconcluded that the explosion cannot be obtained in 1D!
(The exception is an 8.8 M⦿ star; Kitaura+ 06)
Thompson+ 03

Oct 29th 2013

Sumiyoshi+ 05

Supernovae and Gamma-Ray Bursts 2013

5 /21

Neutrino-driven explosion in multi-D simulation
The Astrophysical Journal, 756:84 (22pp), 2012 September 1

The Astrophysical Journal Letters, 767:L6 (7pp), 2013 April 10
Müller, Janka, & Marek

12 ms
12 ms

Bruenn et al.

90 ms

Recently, we have successful exploding models driven by neutrino heating
100
50

100

YS, Kotake, Takiwaki, Whitehouse, Liebendörfer, Sato, PASJ,
0 62, L49 (2010)-150 -100 -50
0
50
400

8 12.5 17

22

100

150

0

-300

-200

-100

0

100

-400

-200

0

200

200

300

comparison
between 1D and 2D
150 ms200 msEntropy [kB-1]
2

400

26 30

200

200

0

Distance from symmetry axis [km]

200

0
200

200
Radial velocity [km s-1]

400

-30000

0

-400

400

30000

-200

0

200

400

300 ms

1000

-600

1000

0

0

500

1000

1000

2000

Müller, Janka, Marek (2012)-1500 -1000

-500

0

500

1000

1500

600 ms

4000

0

1000

0

0

2000

3000

2000 3000

800 ms

6000
3000

4000

Brruenn et al. (2013)-3000 -2000 -1000

2000

600

400
400ms
ms

2000

500

400

6000
-6000 -4000 -2000

0

2000

4000

6000

-9000

-6000 -3000

0

3000

6000

9000

Distance along symmetry axis [km]Oct 29th 2013

Figure 1. Evolution of the entropy (upper half) and radial velocity (lower half) for B12-WH07, with snapshots at tpb = 12, 90, 150, 200, 300, 400, 600, and 800 ms.

Supernovae and Gamma-Ray Bursts 2013

6 /21

The scalein timecaptureand
thetheexpansion
the baryon
supernova
shockwave,
the color
Figure 6. Snapshots of the evolution of model G11, depicting the radial velocity
vr grows(left halfof topanels)entropyofpers (right
half ofbut
panels)115maps
ms,remain constant. The radial velocity portion is omitted for the firstsnapshots.
203 ms, 290 ms, 490 ms, 658 ms, and 920 ms after bounce (from top left totwo
bottomright).

Problems of neutrino-driven explosion
0100
time [ms]

200

300

too small explosion energy (~1049-1050 erg)

explosion simulation of an 11.2 M! progenitor star. Note that the mass-shell spacing outside of the redcomposition interface between the silicon layer and the oxygen-enriched Si shell) is reduced to steps of

2.0

1.5

1.5
1.0
1.0
0.5

0.0
0

when we terminated
all of the models. F
speculate that they
the next 500 ms by
to unquestionably iin the models, a lo
physics would be ne
Our numerical reresults of Mareka relatively early pshows a more clea
models do.

2.5

0.5

100

200

300

200

300

explosion energy [1049 erg]

radius [103 km]

Marek & Janka (2009)

the conditionsof the 11.2 M!
layer increases. At the same
er between theues to be welll, this suggests
he gain layer tom the continued
re 10). Shortly
d a pronounced
ms, it makes
is decline is ainosities at thend the neutron
steep negative
he composition
gen-enriched Si
at the shock (atating timescale
to an evolutionound state, i.e.,

2.0

0.0

4. Summary andtime [ms]Figure 11. Left panel: mean shock radius (arithmetical average over all
lateral directions, dashed line) and maximum and minimum shock positions
as functions of post-bounce time for our two-dimensional explosion simulationof an 11.2 M! progenitor. Right panel: “explosion energy” of the 11.2 M! star,defined as the total energy (internal plus kinetic plus gravitational) of all massin the gain layer with positive radial velocity, as a function of post-bounce time.

Performing 2D cwith spectral neutr
source approximati
expansion of the shsion on the initialdriven outward bymulti-D effects, suc
shown a preponder
with rotation. We hcan become larger f
The conclusion
obtained in this stu(2009), who sugg
impact on the exp
the shock wave on
while the nonrotati
to the short simul
Therefore, because

continuous accretion <=> The remnant is NOT a NS
the absolute value of the total gas energy in the numerator of
Equation (5) goes to zero.
3.4. Explosion Energy
In both our 11.2 M! and 15 M! explosions, the energy ofthe matter in the gain layer with positive radial velocities(“explosion energy”) reaches ∼2.5 × 1049 erg at the end of
the computed evolutions and rises with a very steep gradient(Figures 9 and 11). Therefore, reliable estimates of the final
explosion energy cannot be given at this time. For that to bepossible, the simulations would have to be continued for many
hundred milliseconds more (which is numerically a challenging
task and currently impossible for us with the sophisticatedand computationally expensive neutrino transport and chosen
resolution). This is obvious from the neutrino-driven explosion

Oct 29th 2013

Suwa et al. (2010, 2013)

the panels for
12). Therefore
for some time,
her abundances
he matter swept
nuclei and αh is constructed
a at each radius,
al appears to be(at post-bounce
penetration into
∼ 250 ms p.b.,ening effect on

Fig. 3. Snapshots of the density (left half) and the entropy (right half) for models M13-2D (left panel
shock reaches to 1000 km, corresponding to ! 470 ms after a bounce in both cases.

Fig. 4. Time evolution of the diagnostic energy versus postbounce
time for 2D models with and without rotation.

the 2D models with and without rotation. Although the diagnostic energies depend on the numerical resolutions quantitatively, they show a continuous increase for the rotating models.
The diagnostic energies for the models without rotation, on theother hand, peak at around 180 ms when the neutrino-driven
explosion sets in (see also figure 1), and show a decrease lateron. With values of order 1049 erg it is not yet clear whether
these models will also eventually lead to an explosion.
The reason for the greater explosion energy for models withrotation is due to the bigger mass of the exploding material.
This is because a north–south symmetric (` = 2) explosion can
expel more material than a unipolar explosion can. In fact,
the mass enclosed inside the gain radius is shown to be larger
for the rotating models (e.g., table 1). The explosion energies

Supernovae and Gamma-Ray Bursts 2013

7 /21

The first 3D simulation with neutrino transfer
Takiwaki, Kotake, YS, ApJ, 749, 98 (2012) & recently submitted, arXiv:1308.5755

K computer

T2K-Tsukuba

XT4@NAOJ

320(r)x64(θ)x128(φ)x20(Eν)

Oct 29th 2013

Supernovae and Gamma-Ray Bursts 2013

8 /21

Systematics in supernova simulations
Our Goal: Produce Successful Explosion! of ~1051 erg
08, Nordhaus+ 10, Hanke+ 11,
Dimensionality of hydrodynamics Iwakami+
Takiwaki+ 12, Couch 12, Ott+ 13
General relativity

Liebendörfer+01, Müller+ 12, Kuroda+ 12

Neutrino physics
Scheme to solve Boltzmann equation

Ott+ 08, Shibata+ 11, Sumiyoshi & Yamada 12

Interaction rate

Langanke+ 03, Arcones+ 08, Lentz+ 12

Collective oscillation

Raffelt & Smirnov 07, Duan+ 10,
Dasgupta+ 10

Nuclear equation of state
Initial condition
progenitor structure (mixing, wind...)Lattimer & Swesty 91, H. Shen+ 98, G.
Shen+ 10, Furusawa+ 11, Hempel+ 12
Nomoto & Hashimoto 88, Woosley &Weaver 95, Woosley+ 02, Limongi & Chieffi
06, Woosley & Heger 07, Yoshida+ 12

rotation / magnetic ﬁeld
Oct 29th 2013

Supernovae and Gamma-Ray Bursts 20139 /21

temperature can e.g. be found in Nadyozhin and Yudin
0+ξfRMF(T, nn /ξ, np /ξ) ,(2005); Potekhin et al. (2009); Potekhin and Chabrier(2010). However, we leave this for future studies as the
The first term in Eq. (6) is the summed ideal g
Coulomb energy becomes only important at low temperpression of the nuclei. The Coulomb free energy
atures so that the simplest expression is sufficient for our
nuclei appears in addition. The second line in Eq
purposes.
the direct contribution from the excluded volum(1985)
temperature leads to the population of Hillebrandt
excitedcause&of Wolﬀ
this term,
as long as nuclei are present, tLattimer states&FiniteSwesty(LS)(1991)of the nuclei. Here we use the temperature depenenergy density goes to infinity when approachingdent degeneracy function of Fái and Randrup (1982).Hartree-FockItration density,
because the free volume of nuclei gcalculation
based oncompressible
drop
model
is thesame analyticliquid
expression
as inthe original reference
zero, κ → 0. Thus, nuclei will always disappearof the HS model (Hempel and Schaffner-Bielich 2010),
saturation density is reached. The RMF contribu
G.Shenetal. (2010,2011)
0
nowK=180,
we consider
energies below thevariantsbut
with
220,only
andexcitation
375 MeV
the nucleons fRMF
is weighted with their filling fa
binding energy of the corresponding nucleus, in order toas the free energy is an extensive quantity. If nuc
representthat the2011)
excited states still have to be boundrelativistic
mean
(NL3,
FSUGold)
absent,ξ = ﬁeld
1, andtheory
we get theunmodified
RMF dH.Shen etal.
(1998,(see, e.g., Röpke (1984)). We note that the inclusion oftion, as it should be. The excluded volume correct
excited states up to infinite energies had only aHempel
minor
the al.
nucleirepresents a hard-core repulsion of theet(2012)
relativistic
meanonﬁeld
theory (TM1)
influencethe composition
but would lead to an unat large densities close to saturation density. Instephysically large contribution of the excited states to themodification of the free energy of the unbound nu
relativisticmean
ﬁeld theory
(TM1,
TMA,
energydensitycomponent
and entropy (~2011)at very large temperatures.
including
hyperonis purely
geometric
and just
describes
that the nu
We describe nuclear matter as a chemical mixtureFSUGold)offill only a fraction of the total volume. In this sen
the different nuclear species and nucleons. As we distintwo aforementioned model assumptions for the ex
guish
between nuclei andthe surroundingslope
interacting
nuvolume are essential, as they lead to the desired li
incompressibility
symmetry energy
of symmetry energy
cleons we still have to specify how the system is composed
behavior of the EOS.
K
[MeV]
J
(S)
[MeV]
L
[MeV]
of the different particles. Our thermodynamic model is
2.4. EOS characteristics & constraints
built on two main assumptions: First, we assume for unbound
not allowed to be situatedLS
180,nucleons
220, 375that they are
29.3--Table 2 lists some characteristic saturation pro
inside of nuclei, whereas nuclei are described as uniformof uniform bulk nuclear matter for the three di0
hard spheres at saturation density nB . Second, for nuRMF parameterizations. We also include the L
mass number A ≥36.9
2) we assume that they 111
mustwith the compressibility of K = 180 MeV in the
HShen clei (with281
not overlap with any other baryon in the system (nuclei
The quantities shown in Table 2 correspond to t
or unbound nucleons). Thus we take the volume whichefficients of the following power-series expansion
is available
to be the part of the totalbinding energy per baryon at T = 0 around the s
HW
263 for the nucleons32.9--volume of the system which is not excluded by nuclei.
tion point:
This is described by the filling factor of the nucleons271.5 (NL3)
37.29
(NL3)
118.2 (NL3)
1
1
2
!
GShen
E(x,
β)
=
−E
+
Kx
+
K ! x3 + ...
00
162
230.0 (FSU)ξ = 1 −32.59A (FSU)
60.5 (FSU)
nA,Z /nB ,
(3)
" 18
#
1
A,Z
318 (TMA)
30.7
(TMA)
90 (TMA)+β 2 J + Lx + ... + ... ,Hempel3
32.6we(FSU)
60 (FSU)(here230
and(FSU)in the following,mean A ≥ 2). The
freevolume in which a nucleus can move is the total volume
with x = nB /n0B − 1 denoting the relative deviatio
Oct 29th 2013
10/21
Supernovae and Gamma-Ray Bursts 2013

Finite temperature EOSs

Numerical simulation
400

EOS: LS180, (LS220,) LS375, and Shen
Axisymmetric simulation (ZEUS-2D; Stone & Norman 92)
Hydrodynamics + Neutrino transfer
df
f
+µ
+ µ
cdt
r
= j(1

d ln
3v
+
cdt
cr
f)

(1

f
µ )
+ µ2
µ
2

E2
f+
(1
c(hc)3

f)

d ln
3v
+
cdt
cr

Rf dµ

f

Imcompressibility [MeV]

LS375

v
f
D
cr
E
R(1

350
300

Shen

250

LS220

200
150

LS180
20

f )dµ

25

30

35

40

Symmetry Energy [MeV]

Note: Of course the other parameters
differ as well.

(Lindquist 1966; Castor 1972; Mezzacappa & Bruenn 1993)

Isotropic Diﬀusion Source Approximation (Liebendörfer+ 09)
Ray-by-Ray plus
electron-type neutrino/antineutrino

progenitor: 15 M⦿ (Woosley & Weaver 95)

Oct 29th 2013

Supernovae and Gamma-Ray Bursts 201311/21

Shock radius evolution depending on EOS
YS, Takiwaki, Kotake, Fischer, Liebendörfer, Sato, ApJ 764, 99 (2013)

maximum

15 M⦿ of Woosley& Weaver (1995)

averageminimum

LS180 and LS375 succeed the explosion
Shen EOS fails
Oct 29th 2013

Supernovae and Gamma-Ray Bursts 2013

O’Connor & Ott 10
12/21

6

0

100

200

300

400

Radius of neutron
starTime after Bounce [ms]

500

600

(b)YS,
Mass-weighted
entropy per
gain region.
Takiwaki, Kotake,average
Fischer, Liebendörfer,
Sato,baryon
ApJ 764,in
99 the(2013)
Janka, arXiv:1206.2503

90

70
60
50

80

60

40

20

60

40

20

L&S
H&W
Shen

0
0

100

200

300

time [ms]

400

500

1500

L&S
H&W
Shen

FIG. 4: Two-dimensional SN simulations [121] of an 11.
show cross-sectional entropy distributions at 412 ms after(middle), and at
500 ms p.b. for the Hillebrandt & Wolff Eo
1000
contraction of the PNS (bottom left) and because of weak
L&S
does not yield an explosion within the simulated time as v
shock radius [km]80

Faster contraction isbetter for the
explosion!

80

neutron star radius [km]

LS180
LS375
Shen

neutron star radius [km]Radius of Neutron Star [km]

100

500

H&W
Shen

strong feedback between the EoS, weak interactions, n
In order to achieve a more elaborate treatment of t
0
100
200
300
400
500 trinospheric
600 densities after bounce and to connect sm
00
11 400 3 500
0
100
200
300
400
500
0
100
200ρ < 300
Garching
simulations
employat
10g/cm a Bo
Time after Bounce [ms]
time
[ms]
time [ms]species, and in the non-NSE regime at T ! 5 × 109 K−3(d) Protoneutron star radii, determined at ρ = 1011 g cm
now, a .small reaction network for nuclear burning.
FIG. 4: Two-dimensional SN simulations [121] of an 11.2
M!astar[22] forgravitational
three differentmass
nuclear
EoSs.
With
maximumof 1.83
M!The
forshow cross-sectional entropy distributions at 412 ms after bounce for the LS180-EoS (left), at 586 ms p.b. for thwith PSR J1614-2230. Moreover, an incompressibilit(middle),
and
at
500
ms
p.b.
for
the
Hillebrandt&
Wolff
EoS [122]. The last is the stiffest EoS of the set. It leads t
the models using the LS EOS have a larger anisotropic downflow,
favoredvalue ofheatingK ∼ 240for symmetricnuclearm
contraction of the PNS (bottom left) and because of weaker neutrino
andMeV
less vigoroushydrodynamicndOct
1029thso 2013that the downflow
is yieldrathercloseandtoGamma-Ray
bespherical,
slightly
both
constraints,of shock∼15 km
a 1
13for
/21
does notan explosion
withinthe simulated
time fulfills
as2013
visiblein theevolution ofitstheradiusaverageradius(bott
Supernovae
BurstsProgenitor dependence would be more critical
When we use 11.2 M⦿ as an initial condition

ejectaBoth EOS explode!

NS mass~1.3 MNS
YS, Takiwaki, Kotake, Fischer, Liebendörfer, Sato, ApJ 764, 99 (2013)

Oct 29th 2013

But still the explosion
enegy is ~1050 erg...

Supernovae and Gamma-Ray Bursts 2013

14/21

Systematics in supernova simulations
Our Goal: Produce Successful Explosion! of ~1051 erg
08, Nordhaus+ 10, Hanke+ 11,
Dimensionality of hydrodynamics Iwakami+
Takiwaki+ 12, Couch 12, Ott+ 13
General relativity

Liebendörfer+01, Müller+ 12, Kuroda+ 12

Neutrino physics
Scheme to solve Boltzmann equation

Ott+ 08, Shibata+ 11, Sumiyoshi & Yamada 12

Interaction rate

Langanke+ 03, Arcones+ 08, Lentz+ 12

Collective oscillation

Raffelt & Smirnov 07, Duan+ 10,
Dasgupta+ 10

Nuclear equation of state
Initial condition
progenitor structure (mixing, wind...)Lattimer & Swesty 91, H. Shen+ 98, G.
Shen+ 10, Furusawa+ 11, Hempel+ 12
Nomoto & Hashimoto 88, Woosley &Weaver 95, Woosley+ 02, Limongi & Chieffi
06, Woosley & Heger 07, Yoshida+ 12

rotation / magnetic ﬁeld
Oct 29th 2013

Supernovae and Gamma-Ray Bursts 2013

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NH: Nomoto & Hashimoto (1988)
WW: Woosley & Weaver (1995)
WHW: Woosley, Heger, & Weaver (2002)

Progenitor dependence

YS, Kotake, Takiwaki, Liebendörfer, & Sato (2011)15

10

N13 (NH88)
N15 (NH88)
S15 (WW95)
H15 (WHW02)
H20 (WHW02)
H25 (WHW02)

1014

Density [g cm-3]

1013
1012
1011
1010
109
108
7

10

106

0

0.5

1
1.5
Enclosed Mass [M!]2

2.5

Density proﬁles 100 ms after the bounce
Almost same for M<0.8M⦿

Proﬁle for M>0.8M⦿ reﬂect the initial proﬁleOct 29th 2013

Supernovae and Gamma-Ray Bursts 2013

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Progenitor dependence

NH: Nomoto & Hashimoto (1988)
WW: Woosley & Weaver (1995)
WHW: Woosley, Heger, & Weaver (2002)

13M⦿
NH88

15M⦿
NH88

15M⦿
WW95

15M⦿
WHW02

Oct 29th 2013

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Progenitor dependence

NH: Nomoto & Hashimoto (1988)
WW: Woosley & Weaver (1995)
WHW: Woosley, Heger, & Weaver (2002)

13M⦿
NH88

15M⦿
NH88

Ex
e!

od

e!

od

pl

pl

Ex

l!

i
Fa

Oct 29th 2013

e!

od

pl

Ex
15M⦿
WW95

Supernovae and Gamma-Ray Bursts 2013

15M⦿
WHW02
17/21

Shock evolution in 2D simulation
2D simulation using progenitors from Woosley & Heger (2007)
1000

WH07/s12
WH07/s15
WH07/s20
WH07/s30
WH07/s40
WH07/s55
WH07/s80
WH07/s100

Shock Radius [km]

800
600
400
200
0

0

200

400

600

800

1000

1200

Time after Bounce [ms]Several progenitors lead to shock expansion
No monotonic trend is found
What determines the diﬀerence?
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What makes difference?: Ṁ-Lν curve•

Total Neutrino Luminosity [1052 erg s-1]

Burrows & Goshy 93
1098

explode

7

s12
s15
s20
s30
s40
s50
fail
s55
marginal s80
s100

6
5
4
3
2

0.1

0.2

0.3 0.4 0.5 0.6 0.7 0.8 0.9
Mass accretion rate at 300 km [M! s-1]

1

Fig. 4.— The time evolution of simulated modles in the Ṁ -L
Low M and high Lν are achievedfor Longer
severalsimulations
progenitors,
which produce theplane.
are required!
•explosion
In order to unveil the relationship between the progenitor structure andtrajectories in this plane, more systematic study is necessary...
Oct 29th 2013

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Discussion
O’Connor & Ott 11

Mass
12
15
20
30
40
55
80
100

Explosion / Failure

◯

△
◯
◯
△

Smaller presupernova mass seems to be better for at least neutrinodriven explosion
Model systematics around 15-30 are smaller compared to the othermass range and these progenitors are diﬃcult to explodeOct 29th 2013

Supernovae and Gamma-Ray Bursts 2013

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Summary
For supernova modeling, there are a lot of ingredients topin down the explosion mechanism

We performed multi-dimensional neutrino-radiationhydrodynamic simulations of core-collapse supernovae
The physical parts investigated are
Multi dimensionality [1D<2D>?3D] (YS+ 2010; Takiwaki, Kotake, & YS 2012, 2013)
Eﬀect of neutrino oscillation [potentially strengthen the explosion] (YS+ 2011)
Impacts of nuclear equation of state [“softer” is better] (YS+ 2013)
Dependence of Progenitor structure [under investigation...] (YS+ in progress)

There are still a lot of tasks to do to unveil the explosion
mechanism of core-collapse supernovae...
Oct 29th 2013

Supernovae and Gamma-Ray Bursts 201321/21