中性子星連星の合体に対する電磁波対応天体仏坂健太（東京大学 ビッグバン宇宙国際研究センター）© Tohuku University© Google map~20km~1.4Msun© Tohuku University

1. 重力波の強力な発生源~20km~1.4Msun

KAGRA© Tohuku University

1. 重力波の強力な発生源

Fujimoto, Fukushima, KH, Kyutoku 22

2. 超高密度物質の実験場~20km~1.4Msun© Tohuku University© Gerry Images/iStockphoto

1. 重力波の強力な発生源
2. 超高密度物質の実験場~20km~1.4Msun

3. 金・プラチナ等、重元素の生成現場？

© Tohuku University

1. 重力波の強力な発生源
2. 超高密度物質の実験場~20km~1.4Msun

3. 金・プラチナ等、重元素の生成現場？
https://www.swift.ac.uk/about/grb.php

4. ガンマ線バーストの起源？

© Tohuku University

1. 重力波の強力な発生源？

重力波宇宙論2. 超高密度物質の実験場？
~20km~1.4Msun

3. 金・プラチナ等、重元素の生成現場？
4. ガンマ線バーストの起源？
5. 重力波宇宙論を可能にする？
© Tohuku University

EM Counterpart Discovery Timeline

中性子星合体 GW170817の発見

Gravitational waves(2017 Aug 17.5)
T=0Near Infrared
T = 11h 36m

X-rays
T = 9 days

Off-axis jet?

Delayed Arrival of X-rays and Radio

Gamma-rays
T = 1.7 seconds

OpticalT = 10h 52m

Ultraviolet
T = 15 hours

Radio
T = 16 days

Extremely faint!
(~15 uJy @ 3 GHz)
Abbott et al., ApJL. 848, L12 (2017)

本日の話題•

中性子星合体とは？

•電磁波対応天体

キロノバ•

電磁波対応天体

シンクロトロン残光•展望

連星中性子星とその合体
中性子星

中性子星

重力波
重力波が軌道エネルギーと角運動量を持ち去り合体に至る•

合体までに要する時間 ~ 0.3億年 (at 軌道半径 ~ 太陽半径)※合体時間は軌道半径の４乗に比例•

徐々に軌道半径が減少し、合体に至る。

連星中性子星とその合体
中性子星

GW170817
中性子星

重力波
重力波が軌道エネルギーと角運動量を持ち去り合体に至る

軌道半径~100km

合体~10km

EOS=APR, Mtot = 2.7Msun

EOS=APR, Mtot = 2.9Msun

KH + 13中性子星合体から飛び出す物質1 合体の瞬間
KH, et al 13

私の博士論文のメインテーマ。数値相対論による質量放出の最初の計算。

中性子星合体後に何ができる？

Hypermassive neutron star
(HMNS)
Lifetime ~ 10 ms-1s
Black Hole + disk

(Super)massive neutron star
(SMNS)

Total binary mass

中性子星合体から飛び出す物質2 合体後
14

Dynamical ejecta

FIG. 6. Schematic picture of the ejecta profile for the caseof a sti↵ EOS in which a long-lived massive neutron star isformed as a remnant. The largest anisotropic-shell component(red color) denotes the neutron-rich dynamical ejecta. The
smaller anisotropic-shell component (blue color) denotes theearly viscosity-driven ejecta and long-term viscosity-driven
ejecta from the torus. The polar spheroid component (darkblue color) denotes the viscosity-driven ejecta from the torus
influenced by neutrino irradiation from the massive neutronstar. The “Low Ye ” implies that it contains neutron-rich matter with Ye . 0.2, which contributes to enhancing the opacity through the nucleosynthesis of lanthanide elements. The“Medium Ye ” and “High Ye ” imply that it does not containsuch neutron-rich matter because Ye & 0.25 and Ye & 0.35,
respectively. The filled (purple) circle and neighbouring small

viscosity-driven ejecta is sufficiently high as & 0.03M(i.e., ↵vis is sufficiently
large ⇠ 0.02–0.04),
the electroShibata+17,
Fujibayashi+18
magnetic observations for GW170817 can be naturallyinterpreted.
The
Astrophysical(20pp), 2018 Juneejec10One
unclear pointJournal,
in the 860:64
early viscosity-drivention is that we do not know whether ↵vis is really sufficiently large ⇠ 0.02–0.04 around the central regionof the remnant massive neutron star, i.e., a sufficiently
strong turbulence state is realized or not there, although↵vis = O(0.01) is a reasonable magnitude for turbulent
fluids: Indeed, our latest high-resolution MHD simulation [72] shows that at least for an outer region of theremnant massive neutron star and torus, ↵vis is likely to
be enhanced to ⇠ 0.02. To assess the validity of this
scenario, however, we need to perform a high-resolution
MHD simulation for the merger and post-merger of binary neutron stars, in which several MHD instabilities
such as Kelvin-Helmholtz and magneto-rotational instabilities are well resolved. We note that if the initial torus
mass of the merger remnant is higher (e.g., for the mergerof significant binary mass asymmetry), the ejecta massof ⇠ 0.03M may be achieved for a smaller value of ↵vis .
Thus, the required value for ↵vis may be smaller.
In this section, we have paid particular attention tothe optical-IR counterparts in the relatively early phaseof . 5 days. In the late phase, the e↵ect of the dynamical
mass ejection of low Ye (i.e., of high values of ) should
be visible. The late-time reddening [6, 10, 84] is likely
to be associated with the dynamical ejecta component inour scenario.

合体後、中性子星＋降着円盤、もしくはBH＋降着円盤から質量放出が起こる。アウトフローの性質は、中性子星合体後の進化に強く依存。
IV. DISCUSSION
A.

Perspective for constraining the neutron-star本日の話題•

中性子星合体とは？

•電磁波対応天体

キロノバ•

電磁波対応天体

シンクロトロン残光•

展望

電磁波対応天体 キロノバ
Li & Paczynski 98, Kulkarni 05, Metzger + 10, Barnes & Kasen 13, Tanaka & KH 13

Merger

質量放出

R過程元素合成

放射性加熱
光が抜け始める

晴れ上がり

Time
t=0

0<t<100ms

~<1s

1day

> 10day

10km

10-100km< 0.1Rsun

10AU>100AU中性子星合体のR過程元素合成
原子核

中性子

潮汐破壊

•

中性子星の一部が引き伸ばされ、中性子が原子核に捕獲される非平衡反応(R過程)が起こる。

•

ビッグバン元素合成に類似：ビッグバンでは陽子過剰で光が極めて多い熱い環境という点で異なる。

中性子星合体のR過程元素合成
第 V 章 ビッグバンモデルの提唱

191

αβγ 1948

組成比 (対数)

原子核

中性子

潮汐破壊
質量数

期に，ρn dt の積分値が 5×104 g sec/cm3 に等しい必要があ
る†3 ．

•

中性子星の一部が引き伸ばされ、中性子が原子核に捕獲される非平衡反応(R過程)が起こる。
*4•

ビッグバン元素合成に類似：ビッグバンでは陽子過剰で光が極めて多い熱い環境という点で異し，ρ ≈ 106 /t2 で与えられる．この時間積分は時刻 t = 0 で一方，相対論的な膨張宇宙によると

密度は時間に依存発散するため，元素合成は以下の関係を満たす時刻
t0 に開なる。Alpher,
Bethe, Gamow 1948は純中性子を初期条件に取ったのでR過程を見ていたとも始されたと考える必要がある
言える。
!
∞

t0

(106 /t2 )dt ≈ 5 × 104 .

（2）

中性子星合体のR過程元素合成© Jonus Lippuner

陽子数

R過程とキロノバのエネルギー源

線

定
安

の

核
子

原

程
過
R

中性子数

多くのベータ崩壊チェーン

β-decay

Nuclear γ-rays

-5

10

n

10-6

10

rest framev = 0.3cv = 0.05c

p

GF

-7

10

10-8
10-9
10-10

1

10

100
Energy [keV]

1000

-5 -6
10

ν

10-5
10-6
10-7
-8

rest framev = 0.3cv = 0.05c

10-6
10-8

10-7

10-8-1010

10-9

Excited1105100
Energy [keV]
10

1000
50

10000
100

Energy (keV)

10day, 3Mpc, 0.01Msun

Figure S6: Predicted X-ray flux from atomic line emission of high atomic mass elements. The predicted fl-4rest frame
falls well below the10sensitivity of our NuSTAR data.
v = 0.3cv = 0.05c-5hotons/s/keV/cm2]

hotons/s/keV/cm2]

10rest framev = 0.3cv = 0.05c

NuSTAR upper limits (Epoch 3+4+5)
Nuclear line emission (5 d Mej=0.05 Msun)
Nuclear line emission (5 d Mej=0.03 Msun)
Nuclear line emission (21 d Mej=0.05 Msun)
Nuclear line emission (21 d Mej=0.03 Msun)

10-10
10-12

10000

γ

3day, 3Mpc, 0.01Msun
NuSTAR upper limits (Epoch 2)

10

5day, 3Mpc, 0.01Msun
-4

-4

10-4
Flux (photons cm-2 s-1 keV-1)

10-4

e-

1day, 3Mpc, 0.01Msun
Flux [photons/s/keV/cm2]

Flux [photons/s/keV/cm2]

KH+16

GW170817 (NuStar limit), Evans..KH+17

10

10-6
10-7
-8

たくさんの核種が崩壊するとき

Decay rate

(Metzger et al 10, Goriely et al 11, Roberts et al 11, Korobkin et al 12,
Wanajo et al 2014, Lippuner and Roberts 2015, KH, Sari, Piran 2017)

P Ni
⌧i

1
/ t

Time

キロノバのエネルギー源：ベータ崩壊•

E(t)Q̇(t) / t

dN
1
/
dt
t
•

KH,Sari,Piran 17

たくさんの核種の崩壊:

核の平均寿命と崩壊エネルギーの関係:

⌧ /E

5

Q̇(t) / t

1.2

キロノバのエネルギー源：ベータ崩壊•

E(t)Q̇(t) / t

dN
1
/
dt
t
•

核の平均寿命と崩壊エネルギーの関係:⌧ /E•

KH,Sari,Piran 17

たくさんの核種の崩壊:

5Q̇(t) / t

1.2

ベータ崩壊の物理定数:
Energy: me c2
Time:2⇡ 3 ~7
tF ⌘ G2 m5 c4 ⇡ 9000 seF

イオン当たりの加熱率:

⇣2⇣
⌘ 1.2
⌘ 1.2
1.2
1 m1e c me c tt⇠
1t
Q̇(t)Q̇(t)⇠ hM⇠
day eV/s/ion
i hM
tFi tFtF tF
2

108 x 太陽光度!

at 1 dayキロノバのエネルギー源：ベータ崩壊
Way & Wigner 1948

KH & Nakar 2020

K. %AY AN l3 E- P. WI GNE R

1016

IQ-

«I

10

Heating rate [erg/g/s]

10

IO

«

lO

C-.

+

IF-

14

t-1.3

1012

10

10

108

g'

I

I

.OOI

10

.

I

.I

l

NI

SIONIOI

Q

lO

Q

i~

I6
I deyap

Is

KPdoye ICPOd~a

核分裂後のベータ加熱（原子炉）

10

6

Fj:G. 6. Theoretical and experimental results for-2the radiotivity t seconds after fission.
of P-particles
P(t) = number
emitted per second, C7(t) samefor y-rays' B{t}=energy of Pparticles in Mev per second,F(t) =energy of y-rays in sameunits. 38(t)+I'(t) is total energy(including neutrino energy) lib-

erated.

10

10

0

2

10
10
Time [s]

4

10

6

10

8

R過程後のベータ加熱（キロノバ）

I (I)FSNl g KNO

ベータ加熱率は、組成パターンの詳細には弱くしか依存しない。

キロノバのエネルギー源：ベータ崩壊
Way & Wigner 1948
IQ-

«I

IO

«

lO

C-.

+

Heating rate [erg/g/s]

10

IF-

g'

KH & Nakar 2020

K. %AY AN l3 E- P. WI GNE R

1016
1014
1012
1010
108

I

I-2 .

.OOI
106 .

I

l

I

10

0

SIONIOI

10

NI

Q

lO

10
10
Time [s]核分裂後のベータ加熱（原子炉）
10

2

4

Q

I6

Is

p ICPOd~ai~ 106 KPdoye
108
I deyaFj:G. 6. Theoretical and experimental results for the radiotivity t seconds after fission.
of P-particles
P(t) = number
emitted per second, C7(t) samefor y-rays' B{t}=energy of Pparticles in Mev per second,F(t) =energy of y-rays in sameunits. 38(t)+I'(t) is total energy(including neutrino energy) lib-

R過程後のベータ加熱（キロノバ）
erated.

I (I)FSNl g KNO

ベータ加熱率は、組成パターンの詳細には弱くしか依存しない。

Observed magnitude

21

キロノバ放射予想
The Astrophysical Journal Letters, 778:L16 (5pp), 2013 November 20

光子z band200 Mpc

22
23

4m

Tanaka & KH 2013
8m

24
25

Hotokezaka et al.

Observed magnitude

1m

26
27
0

5

10

15

20

Days after the merger

近赤外の光度曲線

321

H band200 Mpc

22

4m

23

Solar mass fraction

Figure 1.6 Rest-mass density profiles on the meridional plane for the NS–NS (SLy, Mtot = 2.7M! , Q = 1.0) (left) and BH–NS (H4, Q = 3, χ = 0.75) (right) models
10
at 8.8 ms
5 after the onset of the merger. The red arrows show the velocity profiles of the ejecta.
24
10
4(A color
10 version of this figure is available in the online journal.)
103
105

Total

25

space

Observed magnitude

Radiative Transfer Simulations for NS Merger Ejecta

Observed
magnitude
等級

20

I
Table 1
simulation
using SACRA code (Yamamoto et al. 2008).
We
10
II
III
26
Parameters of the Progenitor Models and Their Ejecta Properties
IVfollow10 the dynamical ejecta with the numerical-relativity
simulation10until the head of the ejecta reaches "1000 km (see HoEOS
Type RQχ Mej /10−2 M! #vej $/c
1.35 Mtot /M!
27
10
tokezaka et al. 2013 and Kyutoku et al. 2013 for details). After
15
20
APR4 NS–NS 11.10 2.6–2.9 51.0–1.25 10···
0.01–1.4
0.22–0.27
10that, thedensityandvelocity
structuresoftheejectaare
mod0
5
10 15 20 25 30 35 40 45 50 55 60 65 70 75 80 85 90 95
after ·the
SLy
NS–NS 11.4 2.6–2.8 Days
1.0–1.25· · merger
0.8–2.0
0.20–0.26
Atomic number Z
時間
[日]
eled
assuming
homologous
expansion(Rosswogetal.
2013a).
12.4 2.6–2.8 1.0–1.25 · · ·0.15–0.55
0.22–0.24
Fig. 1.— (Upper) Solar abundance r-process abundance raio in mass fraction (Simmerer et al. 2004). The abundance is normalized with ALF2 NS–NS
X(Ge)= 10 the. (Lower)
The number of bound-boundtransition datafor different elements.
Different colors show
different
ionization states,Forsimulations,weemployapiecewise
polytropic
EOS
with
Fig.
observed
ugrizJHK-band
lightcurves(in AB magn
NS–NS 13.6
2.6–2.8
1.0–1.25 · · ·0.03–0.40
0.18–0.26
from neutral (I) to triply ionized (IV) ions. The atomic data at Z ≤ 30 are taken from Kurucz & Bell (1995) while the data
at Z ≥ 318.—
are H4 Expectedcompiled
using
the
VALD
database(Piskunov
et
al.
1995;
Ryabchikova
et
al.
1997;
Kupka
et
al.
1999,
2000).
It
is
shown
that
there
is
no
which重元素が光を強く吸収し、キロノバは従来の期待よりも暗く赤いと予想。
theionscold
EOSs of neutron-star matter are well fitted (Readto the NSMS1
mergerevent
is setto be1.0–1.25
200 Mpc.
correction0.18–0.20
is taken into
NS–NS
14.4
2.6–2.8· · · K0.06–0.35
data for triply ionized(IV, purple) at Z ≥ 31.
et al. 2009). For systematic studies of the dependence of mass
magnitudes
for wide-field telescopes (5σ with 10 min
exposure).
For optica
APR4 BH–NS 11.1 5.4–10.8 3.0–7.0 0.75
0.05–1.0
0.23–0.27
Forejection
the case of SNe,eachradiation at theplateau phaseto ∼ are100 days).
takenALF2or deduced
from those of PTF (Law et 2.0–4.0
al. 2009), 0.25–0.29
CFHT/Megaca
onthepacket
coldis created
EOSsas aofγ-ray
neutron-starmatter,we(upconsider
D3にして初めて”等級”を学ぶ。
BH–NS 12.4 5.4–10.8 3.0–7.0 0.75
packet. According
to the total energy release by Ni(JHK bands), “4 m” and “space” limits are taken or dedtransfer NIR
decayfive(E )coldand EOSs:
Co decay (E), each(Akmal
packet is desAPR4
et al. 1998) and 2.4.
SLyγ-ray(Douchin& wavelengths
H4
BH–NS 13.6 5.4–10.8 3.0–7.0 0.75
4.0–5.0
0.24–0.29
ignated as that from Ni or Co; a fraction E /E
For the case of SNe, γ-ray transfer isWFIRST
computed. We (Green et al. 2012) and WISH (Yamada et al. 2012), respectively
2001)
as soft
ALF2(Alfordet al.
2005) asa amodis a Haensel
Ni packet while
E /E
is a EOSs,
Co packet.
If
BH–NS 14.4 5.4–10.8 3.0–7.0 0.75
6.5–8.0
0.25–0.30
adopt the grayapproximation
withmass absorptionMS1
Number of lines

4
3
2
1
0

6

56

56

Ni

56

56

56

56

Co

Co

rad

56

Ni

rad

2

−1

ガンマ線バーストとキロノバ
Tanaka & KH 2013, KH+2013
The Astrophysical Journal Letters, 778:L16 (5pp), 2013 November 20

Hotokezaka et al.

光子

The Astrophysical Journal Letters, 778:L16 (5pp), 2013 November 20

ガンマ線バースト 130603B

等級

Magnitude (AB)20
22

SLy(Mej=0.02)
H4(Mej=0.004)r
H

24

Solar mass fraction

Transfer Simulations
for NS Merger
Ejectafor the NS–NS (SLy, M 3 = 2.7M , Q = 1.0) (left) and BH–NS (H4, Q = 3, χ = 0.75) (right) models
Figure 1. Rest-mass Radiative
density profiles
on the meridional
plane
tot
!
at 8.8 ms6 after the onset of the merger. The red arrows show the velocity profiles of the ejecta.
10
526(A color
10 version of this figure is available in the online journal.)
4

10

103

Number of lines

Table 110
simulation
using SACRA code (Yamamoto et al. 2008).
We
Total
28
I
Parameters of the Progenitor Models and Their Ejecta Properties
II
follow10 the dynamical ejecta with the numerical-relativity
simuIII
lation10until the head of the ejecta reaches "1000 km IV(see HoEOS
Type R1.35 Mtot /M!
Qχ Mej /10−2 M! #vej $/c10
tokezaka et al. 2013 and Kyutoku et al. 2013 for details). After
3011.1 2.6–2.9 1.0–1.25 · · · 0.01–1.4 0.22–0.27
APR4 NS–NS
10
that, thedensity and velocity structures of the ejecta are mod0.1 2.6–2.8 1.0–1.251 · · · 0.8–2.0 100.20–0.26
SLy
NS–NS 11.4
100
5
10 homologous
15 20 25 30 35expansion
40 45 50 55(Rosswog
60 65 70 75
90 95
eled assuming
et 80al. 852013a).
ALF2 NS–NS 12.4 2.6–2.8 1.0–1.25 · · ·0.15–0.55
0.22–0.24
Atomic number Z
時間[日]
Rest-frame
days
after
GRB
130603BFor
thesimulations,weemployapiecewise
polytropic
EOS
withFig. 1.— (Upper) Solar abundance r-process abundance raio in mass fraction (Simmerer et al. 2004). The abundance is normalized with H4
NS–NS 13.6 2.6–2.8 1.0–1.25 · · ·0.03–0.40
0.18–0.26
X(Ge) = 10 . (Lower) The number of bound-bound transition data for different elements. Different colors show different ionization states,
EOSs
ofatomic
neutron-star
are wellfittedNS–NS
14.4light2.6–2.8
1.0–1.25· · ·BH–NS
0.06–0.35
0.18–0.20
from which
neutral (I) tothe
triplycold
ionized (IV)ions. The
data at Z ≤ 30 are mattertaken from Kurucz& Bell (1995)
while (Readthe data at Z ≥ 31 are MS1
Figure3. Predictedcurves for
NS–NS andmodels. Leftpanel: NS–NS mocompiled using the VALD database (Piskunov et al. 1995; Ryabchikova et al. 1997; Kupka et al. 1999, 2000). It is shown that there is no
ettriplyal.ionized2009).
For
systematicstudies of the dependence of mass
for the models: SLy (Q = 1.0, Mej = 0.02 M! ), H4 (Q = 1.25, Mej = 4 × 10−3 M! ), r
data forions (IV,purple)at Z ≥ 31.
APR4
BH–NS
11.1
5.4–10.8
3.0–7.0correspond
0.75
0.05–1.0
0.23–0.27
観測値が、キロノバ理論が予想する時刻、明るさ、色と一致。
middle,and lower
curvesfor each modeltothe high-, fiducialand low-heatejection on the cold EOSs of neutron-star matter, we consider
ALF2
BH–NS
12.4
5.4–10.8
3.0–7.0
0.75
2.0–4.0
0.25–0.29
curves show the models MS1 (Mej = 0.07 M! ), H4 (Mej = 0.05 M! ), and APR4 (Mej
For the case of SNe, each packet is created as a γ-ray(up to ∼ 100 days).
fiveAccording
cold EOSs:
APR4(Akmalet al.radiation
1998)at the
andplateau
SLyphase(Douchin&packet.
to the totalenergy release
by Ni
H4The thin
BH–NS
13.6
0.75light4.0–5.0
0.24–0.29
and thicklines5.4–10.8
denote the 3.0–7.0
r and H-band
curves. Here weset (Q, χ ) = (3, 02.4. γ-ray transfer
decayHaensel(E ) and 2001)
Co decayas(E soft), eachpacketisdesEOSs,
ALF2(Alford
et
al.
2005)
as
a
modMS114.4 of5.4–10.8
3.0–7.0
curvesBH–NS(dashed lines)
the afterglow
model0.75
of GRB6.5–8.0
130603B in r0.25–0.30
and H-band are plo
ignated as that from Ni or Co; a fraction E /E
5
4
3
2
1
0

6

56

Ni

56

56

56

Co

Ni

rad

For the case of SNe, γ-ray transfer is computed. We

ガンマ線バーストとキロノバ
Tanaka & KH 2013, KH+2013
The Astrophysical Journal Letters, 778:L16 (5pp), 2013 November 20

Hotokezaka et al.

光子観測が発表された日

The Astrophysical Journal Letters, 778:L16 (5pp), 2013 November 20

ガンマ線バースト 130603B

20

等級(悔しそうに)Magnitude (AB)

当たったな。。。

22

SLy(Mej=0.02)
H4(Mej=0.004)r
H

24

Solar mass fraction

Transfer Simulations
for NS Merger
Ejectafor the NS–NS (SLy, M 3 = 2.7M , Q = 1.0) (left) and BH–NS (H4, Q = 3, χ = 0.75) (right) models
Figure 1. Rest-mass Radiative
density profiles
on the meridional
plane
tot
!
at 8.8 ms6 after the onset of the merger. The red arrows show the velocity profiles of the ejecta.
10
526(A color
10 version of this figure is available in the online journal.)

中村さん、当時、10研究室の教授10
simulation
using SACRA code (Yamamoto et al. 2008).
We
Total
4

10

3
5

Table 128
Parameters of the Progenitor Models and Their Ejecta Properties

Number of lines

follow10 the dynamical ejecta with the numerical-relativity simulation10until the head of the ejecta reaches "1000 km (see HoEOS
Type R1.35 Mtot /M!
Qχ Mej /10−2 M! #vej $/c10
tokezaka et al. 2013 and Kyutoku et al. 2013 for details). After
3011.1 2.6–2.9 1.0–1.25 · · · 0.01–1.4 0.22–0.27
APR4 NS–NS
10
that, thedensity and velocity structures of the ejecta are mod0.1 2.6–2.8 1.0–1.251 · · · 0.8–2.0 100.20–0.26
SLy
NS–NS 11.4
100
5
10 homologous
15 20 25 30 35expansion
40 45 50 55(Rosswog
60 65 70 75
90 95
eled assuming
et 80al. 852013a).
ALF2 NS–NS 12.4 2.6–2.8 1.0–1.25 · · ·0.15–0.55
0.22–0.24
Atomic number Z
時間[日]
Rest-frame
days
after
GRB
130603BFor
thesimulations,weemployapiecewise
polytropic
EOS
withFig. 1.— (Upper) Solar abundance r-process abundance raio in mass fraction (Simmerer et al. 2004). The abundance is normalized with H4
NS–NS 13.6 2.6–2.8 1.0–1.25 · · ·0.03–0.40
0.18–0.26
X(Ge) = 10 . (Lower) The number of bound-bound transition data for different elements. Different colors show different ionization states,
EOSs
ofatomic
neutron-star
are wellfittedNS–NS
14.4light2.6–2.8
1.0–1.25· · ·BH–NS
0.06–0.35
0.18–0.20
from which
neutral (I) tothe
triplycold
ionized (IV)ions. The
data at Z ≤ 30 are mattertaken from Kurucz& Bell (1995)
while (Readthe data at Z ≥ 31 are MS1
Figure3. Predictedcurves for
NS–NS andmodels. Leftpanel: NS–NS mocompiled using the VALD database (Piskunov et al. 1995; Ryabchikova et al. 1997; Kupka et al. 1999, 2000). It is shown that there is no
ettriplyal.ionized2009).
For
systematicstudies of the dependence of mass
for the models: SLy (Q = 1.0, Mej = 0.02 M! ), H4 (Q = 1.25, Mej = 4 × 10−3 M! ), r
data forions (IV,purple)at Z ≥ 31.
APR4
BH–NS
11.1
5.4–10.8
3.0–7.0correspond
0.75
0.05–1.0
0.23–0.27
観測値が、キロノバ理論が予想する時刻、明るさ、色と一致。
middle,and lower
curvesfor each modeltothe high-, fiducialand low-heatejection on the cold EOSs of neutron-star matter, we consider
ALF2
BH–NS
12.4
5.4–10.8
3.0–7.0
0.75
2.0–4.0
0.25–0.29
curves show the models MS1 (Mej = 0.07 M! ), H4 (Mej = 0.05 M! ), and APR4 (Mej
For the case of SNe, each packet is created as a γ-ray(up to ∼ 100 days).
fiveAccording
cold EOSs:
APR4(Akmalet al.radiation
1998)at the
andplateau
SLyphase(Douchin&packet.
to the totalenergy release
by Ni
H4The thin
BH–NS
13.6
0.75light4.0–5.0
0.24–0.29
and thicklines5.4–10.8
denote the 3.0–7.0
r and H-band
curves. Here weset (Q, χ ) = (3, 02.4. γ-ray transfer
decayHaensel(E ) and 2001)
Co decayas(E soft), eachpacketisdesEOSs,
ALF2(Alford
et
al.
2005)
as
a
modMS114.4 of5.4–10.8
3.0–7.0
curvesBH–NS(dashed lines)
the afterglow
model0.75
of GRB6.5–8.0
130603B in r0.25–0.30
and H-band are plo
ignated as that from Ni or Co; a fraction E /E
I
II
III
IV

4
3
2
1
0

6

56

Ni

56

56

56

Co

Ni

rad

For the case of SNe, γ-ray transfer is computed. We

ガンマ線バーストとキロノバ
Tanaka & KH 2013, KH+2013
The Astrophysical Journal Letters, 778:L16 (5pp), 2013 November 20

Hotokezaka et al.

光子

翌日

The Astrophysical Journal Letters, 778:L16 (5pp), 2013 November 20

ガンマ線バースト 130603B

20

等級

Magnitude (AB)

キロノバが暗いことの証拠や、こんなん見つけられへんから違う対応天体考えなあかんわ！
Radiative Transfer Simulations for NS Merger Ejecta

22

SLy(Mej=0.02)
H4(Mej=0.004)r
H

24

Solar mass fraction

3 = 2.7M , Q = 1.0) (left) and BH–NS (H4, Q = 3, χ = 0.75) (right) models
Figure 1. Rest-mass density profiles on the meridional plane for the NS–NS (SLy, Mtot!
at 8.8 ms6 after the onset of the merger. The(嬉しそうに)
red arrows show the velocity profiles of the ejecta.
10
526(A color
10 version of this figure is available in the online journal.)

中村さん、当時、10研究室の教授10
simulation
using SACRA code (Yamamoto et al. 2008).
We
Total
4

10

3
5

Table 128
Parameters of the Progenitor Models and Their Ejecta Properties

Number of lines

follow10 the dynamical ejecta with the numerical-relativity simulation10until the head of the100%研究者として学生と接する姿勢
ejecta reaches "1000 km (see HoEOS
Type R1.35 Mtot /M!
Qχ Mej /10−2 M! #vej $/c10
tokezaka et al. 2013 and Kyutoku et al. 2013 for details). After
3011.1 2.6–2.9 1.0–1.25 · · · 0.01–1.4 0.22–0.27
APR4 NS–NS
10に感銘を受ける。
that, thedensity and velocity structuresof the ejecta are mod0.1 2.6–2.8 1.0–1.251 · · · 0.8–2.0 100.20–0.26
SLy
NS–NS 11.4
100
5
10 homologous
15 20 25 30 35expansion
40 45 50 55(Rosswog
60 65 70 75
90 95
eled assuming
et 80al. 852013a).
ALF2 NS–NS 12.4 2.6–2.8 1.0–1.25 · · ·0.15–0.55
0.22–0.24
Atomic number Z
時間[日]
Rest-frame
days
after
GRB
130603BFor
thesimulations,weemployapiecewise
polytropic
EOS
withFig. 1.— (Upper) Solar abundance r-process abundance raio in mass fraction (Simmerer et al. 2004). The abundance is normalized with H4
NS–NS 13.6 2.6–2.8 1.0–1.25 · · ·0.03–0.40
0.18–0.26
X(Ge) = 10 . (Lower) The number of bound-bound transition data for different elements. Different colors show different ionization states,
EOSs
ofatomic
neutron-star
are wellfittedNS–NS
14.4light2.6–2.8
1.0–1.25· · ·BH–NS
0.06–0.35
0.18–0.20
from which
neutral (I) tothe
triplycold
ionized (IV)ions. The
data at Z ≤ 30 are mattertaken from Kurucz& Bell (1995)
while (Readthe data at Z ≥ 31 are MS1
Figure3. Predictedcurves for
NS–NS andmodels. Leftpanel: NS–NS mocompiled using the VALD database (Piskunov et al. 1995; Ryabchikova et al. 1997; Kupka et al. 1999, 2000). It is shown that there is no
ettriplyal.ionized2009).
For
systematicstudies of the dependence of mass
for the models: SLy (Q = 1.0, Mej = 0.02 M! ), H4 (Q = 1.25, Mej = 4 × 10−3 M! ), r
data forions (IV,purple)at Z ≥ 31.
APR4
BH–NS
11.1
5.4–10.8
3.0–7.0correspond
0.75
0.05–1.0
0.23–0.27
観測値が、キロノバ理論が予想する時刻、明るさ、色と一致。
middle,and lower
curvesfor each modeltothe high-, fiducialand low-heatejection on the cold EOSs of neutron-star matter, we consider
ALF2
BH–NS
12.4
5.4–10.8
3.0–7.0
0.75
2.0–4.0
0.25–0.29
curves show the models MS1 (Mej = 0.07 M! ), H4 (Mej = 0.05 M! ), and APR4 (Mej
For the case of SNe, each packet is created as a γ-ray(up to ∼ 100 days).
fiveAccording
cold EOSs:
APR4(Akmalet al.radiation
1998)at the
andplateau
SLyphase(Douchin&packet.
to the totalenergy release
by Ni
H4The thin
BH–NS
13.6
0.75light4.0–5.0
0.24–0.29
and thicklines5.4–10.8
denote the 3.0–7.0
r and H-band
curves. Here weset (Q, χ ) = (3, 02.4. γ-ray transfer
decayHaensel(E ) and 2001)
Co decayas(E soft), eachpacketisdesEOSs,
ALF2(Alford
et
al.
2005)
as
a
modMS114.4 of5.4–10.8
3.0–7.0
curvesBH–NS(dashed lines)
the afterglow
model0.75
of GRB6.5–8.0
130603B in r0.25–0.30
and H-band are plo
ignated as that from Ni or Co; a fraction E /E
I
II
III
IV

4
3
2
1
0

6

56

Ni

56

56

56

Co

Ni

rad

For the case of SNe, γ-ray transfer is computed. We

キロノバ

GW170817© ESO

キロノバの光度曲線（GW170817）

͔

͕

‫ݱ‬

Λ

‫ى‬
ベータ崩壊加熱に沿うように減光

ࣼ
下限

ߟ

δ

Δ

キロノバのスペクトルSr lines in
the kilonova spectrum
60
Sr ii

F [10 17 erg s 1 cm 2 Å ] + offset

50

Sr ii

Blackbody + Sr ii
Blackbodyblueshiftdecades later30 . However, no spectroscopic confirmation of the composition of neutron stars has ever been made. The identification here ofan element that could only have been synthesised so quickly under an
extreme neutron flux, provides the first direct spectroscopic evidence
that neutron stars comprise neutron-rich matter.

X-shooter spectra are explained byblackbody + Sr II lines.

1.5日のスペクトルは、綺麗な黒体放射408に吸収線が載っている。Tanaka & KH 13
Tanaka & Hotokezaka

Log flux (Fλ) + constant1

1. Burbidge, E. M., Burbidge, G. R., Fowler, W. A. & Hoyle, F. Synthesis of
6
the Elements
in Stars. Rev. Mod. Phys. 29, 547–650 (1957).
2. Siegel, D. M., Barnes, J. & Metzger, B. D. Collapsars1.5
as daysa major source
1.5 daysof r-process
5 elements. Nature 569, 241–244 (2019). 5.0 days
30 offsets
3. Lattimer, J. M., Mackie, F., Ravenhall, D. G. & Schramm,
D. N. The
10.0 days
decompression of cold neutron star matter. Astrophys. J. 213, 225–233
4
2.5 days(1977).
4. Eichler, D., Livio, M., Piran, T. & Schramm, D. N. Nucleosynthesis, neu20
3 and gamma-rays from coalescing neutron stars. Nature 340,
trino bursts
3.5 days
126–128 (1989).
5. Freiburghaus,C., Rosswog, S. & Thielemann, F.-K. R-Process in Neutron
2
Star Mergers. Astrophys. J. Lett. 525, L121–L124 (1999).
10
6. Ji, A. P., Frebel, A., Simon, J. D. & Chiti, A. Complete Element Abun4.5 days
1 Nine Stars in the r-process Galaxy Reticulum II. Astrophys. J.
dances of
10000
15000
20000
830, 93 (2016). 5000
7. Metzger, B. D. et al. Electromagnetic
counterpartsof compact object
Wavelength(A)mergers powered by the radioactive decay of r-process nuclei. Mon. Not.
3
1.5 daysFig. 6.—UVOIR spectra
model NSM-all
at 2650–2662
t = 1.5, 5.0, and10.0 days after the merger. The spectra are almost
R.ofAstron.
Soc. 406,(2010).
-3
wavelengths, there are possible absorption troughs, which result from Y i, Y ii, and Lu i in our simulations. Howeve
8. Barnes, J. & Kasen, D. Effect of a High Opacity on the Light Curves ofbe result2.5from
daysthe incompleteness of the line list in the NIR wavelengths.
3Radioactively Powered Transients from Compact Object Mergers. Astro-3• phys. J. 775, 18 (2013).
3
3.5 days
9. Tanvir, N. R. et al. A ‘kilonova’ associatedthe short-duration-ray of GW sof withEM radiation
to the analysis-3
1042
GRB 130603B. Nature 500, 547–549(2013).
able
to pin down the masses of two NS•1.2 burst
3+
1.5
4.5 days
NSM-all
10. Abbott, B. P., Abbott, R. & et al. GW170817:
Observationof Gravitationalof theEOSs moreaccurately. Note th-3
1.3 Waves+ 1.4 from a Binary
APR4Neutron(soft) Star Inspiral. Physical Review Letters 119,
simulations, the heating rate per mass•
H4 (stiff)
4000
5500
7500
10500
14500
20000
161101 (2017).
understand the connection between th
11. Pian, E. et al. Spectroscopic identification
nucleosynthesis
of ofther-process
NS merger
and expected emissio
Observed wavelength [Å]
1041
in a double neutron-star merger. Nature 551,
67–70 calculations(2017).
osynthesis
are necessary.
12. Smartt, S. J. et al. A kilonova as the electromagnetic counterpart to a
Figure 4 | Spectral series of AT2017gfo 1.5–4.5 days after the merger. Data
gravitational-wave source. Nature 551, 75–79 (2017).
6. IMPLICATIONS FOR OBSE
are shown in grey and have been smoothed slightly. A model (solid red lines) 13. Baade, W. & Zwicky, F. Cosmic Rays from Super-novae. Proceedings of

2.5日以降も同じ波長付近に吸収・放射線が見えるSr II＝＞膨張物質による吸収放射の典型的

1

Residual F [10 17 erg s 1 cm 2 Å ]なプロファイル（P-Cygni）

これらの強いラインはSr II と考えられ

Watson et al 19

inosity (erg s-1)

る（Heの可能性もあり）。

Sr lines are expected to be very strong.
The early blackbody is surprising.
Heavy elements may be absent in theouter part of the ejecta.

キロノバ後期スペクトルと元素の起源
GW170817 spectrum rescaled to 100 Mpc (KH, et al 2022)•

10日の近赤外ピーク＝＞Te (Z=52)?

•

40日の赤外観測=> W (Z=74)?

•

JWST(感度~0.1μJy)で見れば、綺麗なスペクトルが取れると期待。

GW170817からの元素の起源への示唆

Watson+ 19, Perego+21, Gillandars+22, KH+22, Domoto+22, Tarumi, KH+ in prep.

本日の話題•

中性子星合体とは？

•電磁波対応天体

キロノバ•

電磁波対応天体

シンクロトロン残光

•

展望

1432
K. Hotokezaka and T. Piranシンクロトロン電波残光
Nakar & Piran 2011, KH & Piran 2015

ジェット

コクーン
Dynamical

膨張速度

ジェット

Fast tailof ejecta

Fast
Fast tail of ejecta

ーン
ク
コ

Brigh

Dynamic
Wind

Wind
運動エネルギー

Figure 1. Left-hand panel: the kinetic energy and the radial component ofdeceleration time-scales of equation (2) assuming an external density of 1 cm
• 中性子星合体のアウトフローは様々な速度・エネルギーを持つ。
schematic picture of the morphology of the different components of the ejecta merger simulation (Hotokezaka et al. 2013a). Other components are added s• アウトフローが星間物質と衝突し衝撃波加速が起こりシンクロトロン放射を生成。
and the l.h.s.

明るさ:

Fpeak / E 11/4 n7/8

（電波にピーク）and β is a velocity in units of the speed of light c. Also shownin the figure are膨張速度
the deceleration
time-scales due to the interaction
運動エネルギー
星間物質の密度with the ISM, which are discussed later. This time-scale gives thethis the case of fast cooling. The flux
iven byK. Hotokezaka and T. Piranシンクロトロン電波残光

1432
Fn,max ,nc 1 n,2Fn,max ,nm 1 n 1 nc ,
(7)
1/2
2p/2
& Piran(n/nm )Nakar
Fn,max, n 12011,
nm , KH & Piran 2015

ジェット

Fast

膨張速度

Fn,max { Ne Pn,max /4pD 2 is the observed
ジェット
from
the tailsource.
Fast
thoseof electronsejecta with ge 1 gc can cool.
ng, because the electrons with ge ª gm,
f the population, do not cool within a

コクーン

Dynamical,
nm 1 n,Fn,max ,nc 1 n 1 nm ,
(8)
Wind
1)/2
(n/nc)2p/2 Fn,max , n 1 nc .

Fast tail of ejecta

ーン
ク
コ

Brigh

Dynamic

加速電子からシンクロトロン予想スペクトル
Wind

,max
21)/2

Sari et al 1998

運動エネルギー

corresponding to fast and slow coolingFig.1.1.—Synchrotron
spectrumof aenergy
relativistic
shock
with
a power-law
FigureLeft-hand panel:
the kinetic
and theradialcomponentofa and 1b. The low-energy part of these!
electron distribution.
cooling,
which is expected
at early
times
deceleration
time-scales(a)ofFast
equation(2) assuming
an external
densityof 1(tcmt0). The spectrum consists of four segments, identified as A, B, C, and D. Self
upport• even
within the GRB itself (Co中性子星合体のアウトフローは様々な速度・エネルギーを持つ。
schematic
picture of the morphology of the different components of the ejectabsorption
is important below na . The frequencies, nm , nc , and na, decrease withition •to the various power-law regimesamergersimulation
et al. 2013a).
Other
components
areadiabatic
added s
アウトフローが星間物質と衝突し衝撃波加速が起こりシンクロトロン放射を生成。
time as indicated; (Hotokezaka
the scalings above
the arrows
correspondto an
bsorption causes a steep cutoff of theandthe l.h.s.and the scalings below, in square brackets, correspond to a fully
evolution,
ncies (Katz 1994; Waxman 1997b; Katz
radiative evolution.
(b)7/8Slow cooling, which is expected at late times (t 1 t0)
11/4
ompleteness, we 明るさ:
show this regimepeakinThe evolution is always adiabatic. The
four segments are identified as E, F
（電波にピーク）G, andgnore it for the rest of this Letter sinceand
β isH.a velocity in units of the speed of light c. Also shownr the optical or the X-ray radiation inin the figure are膨張速度
the deceleration
time-scales due to the interaction
運動エネルギー
星間物質の密度we thewrite
M whichin terms
the initialenergyof the explosion
d.
with
ISM,
are ofdiscussed
later.
This time-scale
gives thevia

F

/E

n

2015年当時、電波残光は対応天体として劣勢
Nakar & Piran (PD時代のボス)

様々な否定的な意見
1. 合体は薄い密度の環境
2. エジェクタの速度が遅い
3. ピーク時間が遅い
4. 母銀河と分離できない
5. AGNと超新星と間違える
6. …..

2015年当時、電波残光は対応天体として劣勢
Nakar & Piran (PD時代のボス)

様々な否定的な意見
1. 合体は薄い密度の環境
2. エジェクタの速度が遅い
3. ピーク時間が遅い
4. 母銀河と分離できない
5. AGNと超新星と間違える
6. …..

全ての否定的な意見に反論せよ
ほぼ喧嘩に見えるような論争によってサイエンスを進展させる姿勢を学ぶ。

2015年当時、電波残光は対応天体として劣勢
The Astrophysical Journal, 831:190 (21pp), 2016 November 10

Nakar & Piran (PD時代のボス)

様々な否定的な意見
1. 合体は薄い密度の環境
2. エジェクタの速度が遅い
3. ピーク時間が遅い
4. 母銀河と分離できない
5. AGNと超新星と間違える
6. …..

KH et al 2016

G. Hallinan

S. Nissanke残光 in GW170817: 美しい冪則
GW170817:
Afterglow spectrum
光度曲線 (Hallinan+17, Mooley+18,
Makhathini+2020)
スペクトル

37.6

27.5

Electron distribution:

37.4
•
28
Figure 1. Upper panel: The panchromatic (radio, optical and X-ray) afterglow light curve of GW170817, color coded according

dN
d e /

2.2

to the observing frequency, up to 940 days post-merger (all data points have 1 errorbars as presented in Table 1; upper limits
are not shown here) using the uniform dataset presented in this work. The light curve is scaled to 3 GHz using the best-fit
37.2
spectral index ( 0.569) derived
28.5 from the MCMC power-law fitting (see §3). Lower panel: The averaged (using moving average;
t/t = 1/17 where t is the width of the kernel and t is the time after merger) light curve (blue data points) shows a generaltrend consistent with power-law rise and decline. In grey are the same data points as shown in the upper panel.

e37

on the pn detector were 29
used to determine periods ofing to Chandra results from similar epochs. Backgroundhigh background
data,
which
were
excludedfrom
ourevents Merger
were extracted
from nearby, source-free circular
m
=
m
m
Time
after
(days)•
1
old in onlynew
analysis. Significant background flaring resulted
regions of 5500 in radius.
0 for obsIDs 0811210101
e were26 and 48 ks filtered exposures
For both GW170817 and NGC 4993, the spectraand 0811212701,
respectively.
Sourceeventswere
thengroupedwithaminimum
of
1
count
per
bin
with
the
•
-1
extracted from a circular region1 with radius 500 centeredheasoft 2tool grppha. We used the X-ray spectral
3
10
10
10
on GW170817. While this is less than 50% of the EPICfitting package xspec v12.10.1 to fit the data. For
Mergerpn encircled energy, it was necessary to use a small re-Time after
GW170817
we fit(days)
the data with an absorbed powerlaw model (tbabs*ztbabs*powerlaw), where tbabs is agion in order to exclude emission from the AGN, which
is only 10 arcsec from the source. We find this does
neutral absorbing column attributed to our own Galaxy,• any large systematic e↵ect when comparnot introduce
fixed at 8.9 ⇥ 1020 cm 2 (Dickey & Lockman 1990), and

光度曲線・スペクトル、共に冪則を示す。
m

1010

x10 7
x10 6

108

x10 5

~5 orders of magnitude in γ

10

6

104
10

よそ8桁の周波数帯に渡って単一べき則。
• Fermi acceleration works well

10-4

GW170817 spanning ⇡ 110.5 - 584.1 days (green points; observerて、電子加速が起こっていることがよくわ
frame); downwards triangles denote 3 upper limits. The upper
limit at ⇡ 584.1
d is measured
from thepowers
median-subtracted
かる。
Kinetic
energytheimage,while all other data points are measured from HOTPANTS residual
images. Alsoafterglow.
shown are a structured jet model and the range oflight curves describing the top 5% of models (black solid and dotdashed lines), and a quasi-spherical outflow model (dotted line; Wu

x10 3

110
d
165
d
172
d
209
d
218
d
297
d

x10 2
x10

100

特に、スペクトルは電波からX線まで、おアウトフローと星間物質の衝撃波によっ
Figure 2. Top: HST/F606W light curve of the afterglow of

x10 4

2

10-2

108

5

X-ray

x10 8

OpticalA single power law: ~ν-0.6
38=> synchrotron radiation(optically thin)
37.8•

Afterglow (F606W)
Structured Jet (WM18)
Quasi-Spherical(WM18)

Flux Density ( Jy)

27

150日

-1

Optical AB Magnitude

26.5

GW170817 HST

log(L) (erg s )26

Fong+19, also Margutti+18
RadioSpectrum

1010

1012

1014

1016

328
d
362
d
584
d

1018

Frequency (Hz)Figure 3. Broad-band SED of the afterglow of GW170817 at nineepochs of our HST observations, spanning ⇡ 110 - 584 days; fluxes
are scaled for clarity. The HST photometry in this paper (green
Hallinan+17,
circles),
radio afterglowMargutti+17,18,(red squares; Margutti Troja+17,19,
et al. (2018); Mooley
et al. Haggard+17,(2018c); Dobie et al.
(2018); Mooley et al. (2018b); Alexander
Ruan+17,Lyman+18,Mooley+18
et al. (2018); Troja et al. (2018c), Hajela et al. in prep.), and X-ray

電波干渉計を使って見るジェット
Breaking the degeneracy: VLBI
Mooley…KH (2018)

Two observations with the HSA(75 d and 230 d post-merger)

電波干渉計を使って見るジェット
Breaking the degeneracy: VLBI

K. P. Mooley & A. T. Deller

Mooley…KH (2018)

Two observations with the HSA(75 d and 230 d post-merger)

超光速ジェット

in GW170817

VLBI resolve the motion of the radio source Mooley…KH (2018)
1, 電波源が動いている
2.7 ミリ秒角 in 155 日
Day 240

Day 75=> ~ 0.5 パーセク (at 40Mpc)

見かけの速度（c）

app = 4.1 ± 0.4 at 41 M

2,電波源は点源と無矛盾

小さい開きを持つ相対論的なジェット

Figure 1: Proper motion of the radio counterpart of GW170817. The centroid offset posi-

ions (shown by 1 errorbars) and 3 -12 contours of the radio source detected 75 d (black)Movie© Ore Gottlieb & Ehud Nakar

ジェットの見込み角の情報=> H0

e Methods)15 . Figure 2 depicts the poste-

KH+19

the GW-only analysis15, 27 . The constraint+4.5
s/Mpc, to 68.1 4.3 km/s/Mpc (median and

3-4% ジェットのモデルの不定性

今後の展望•LIGO/Virgo/KAGRA O4が今年から始まる。感度向上により、検出頻度が上がる。ただし、距離が遠くなるので、電磁波対応天体探査はより難しくなる。

数年•

高感度宇宙望遠鏡JWSTにより、キロノバスペクトルの大部分をhigh S/Nでカバー。元素の特定など。

~10年

•

中性子星合体の多様性や、ブラックホール・中性子星合体の対応天体。

•

イベント数が50ほどになれば、ハッブル定数の測定精度がIa超新星を超える。

•

次世代の地上干渉計へ。高赤方偏移の重力波宇宙論。重力波＋ガンマ線バース
ト。近傍イベントでは中性子星の状態方程式に対して極めて強い示唆。

~15年•

宇宙干渉計。地上からは見えない重力波ソースが見え、質的な変化が起こる。

~100年 • 系内超新星爆発からのニュートリノ・重力波・電磁波観測