r-­‐process	
  in	
  neutron	
  star	  mergers
Shinya	  Wanajo	
  (RIKEN	  iTHES)	
  

	
  	
  
Y.	
  Sekiguchi	  (YITP),	
  N.	
  Nishimura	
  (Keele	  Univ.),	
  K.	
  Kiuchi,	
  K.	
  Kyutoku,	
  M.	
  Shibata	
  (YITP)	
  	
  iTHES-­‐Kavli	
  IPMU-­‐RESCEU	
  Joint	
  MeeIng	
  July	
  7-­‐8,	
  2014,	
  RIKEN	
  Nishina	
  Hall	  contents	

1. overview	  2. mergers	
  with	  neutrinos	  3. r-­‐process	
  novae	
  (kilonovae)	
  

iTHES-­‐IPMU-­‐RESCEU	

Wanajo	2	1. overview	  iTHES-­‐IPMU-­‐RESCEU	

Wanajo	3	origin	
  of	
  gold	  (r-­‐process	
  elements)	
  is	
  sIll	  unknown…	
  iTHES-­‐IPMU-­‐RESCEU	

Wanajo	4	popular	
  r-­‐process	
  scenarios	

core-­‐collapse	
  supernovae	  (since	  Burbidge+1957;	
  	
  
	
   	
  	
  	
  	
  	
  Cameron	
  1957)	
  v n-­‐rich	
  ejecta	
  nearby	
  proto-­‐NS	
  v not	  promising	
  according	
  to	
  recent	  studies	  iTHES-­‐IPMU-­‐RESCEU	

Wanajo	neutron-­‐star	
  mergers	
  (since	  LaWmer+1974;	
   	
   	
  
	
   	
   	
  	
  	
  	
  	
  Symbalisty+1982)	
  v n-­‐rich	
  ejecta	
  from	  coalescing	
  NS-­‐NS	
  or	
  BH-­‐NS	  	  v few	
  nucleosynthesis	
  studies	
  
5	

SN	
  ejecta:	
  not	  so	
  neutron-­‐rich…	

PHYSICAL REVIEW C 86, 065803 (2012)

IUM MODIFICATION OF THE CHARGED-CURRENT . . .

εYe ~ %1+( ,	  $ Lν e εν e + 2Δ '	
  Δ = M −III. PROTO-NEUTRONSTARM EVOLUTION≈ 1.29 MeVnpo illustrate the effect of the correct inclusion of meaneffects in charged-current interaction rates, as well as thertance of thev
nuclearenergy,five
PNS coolingfor	
  symmetryYe	
  <	
  0.5	
  (i.e.,	
  n-­‐rich)	
  els are described here. The models have been evolved
g the multigroup, multiflavor,
general relativistic variableν e − ν e > 4Δ ~ 5 MeV
ngton factor code described in Ref. [13], which follows
ontraction and neutrino
if Lν e losses≈ Lν eof a PNS over the firstseconds of its life. These start from the same post core
iTHES-­‐IPMU-­‐RESCEU	
ce model considered
in Ref. [13] and follow densities
9−3
n to about 10 g cm . Therefore, they do not simulate theW itself but they do encompass the full neutrino decoupling

ε

ε

νe

~ 6 MeV	ννe<εν> (MeV)

ν

16

12

no	
  correcIon	

8
(a)
0.6
0.55
Ye,NDW

While collisional broadening tends to increase both νe andoss sections,vRPAdecrease thebνy	
  e cross secYecorrelations
	
  is	
  determined	
  and enhance the cross section for ν̄e . Given the simplicity
ur model for the p-h interaction, these−results only serven →ofpthe+ erole of correlations.
pture the qualitative
aspectse+
nonetheless	
   demonstrate that changes+expected are small
p →duen +toea proper treatment
pared to corrections
arisinge+	
  mean field effects
in the reaction kinematics. Hence,e following discussion of PNS evolution and neutrinov theseequilibrium	
  
alue	
  
is	
   andra, we set aside
effects due tov
RPAcorrelations
sional broadening, and calculate the neutrino interactions
	
   field energy shifts calculated−1 asincluding the mean#&Lν e ν e − 2Δ
ribed in Ref. [1].	
  

no

n	
o
I
c
e
	
  corrYe > 0.4 !! 	

0.5
0.45
(b)
0.4

0

2

4

Time (s)6810Roberts+2012
FIG. 5. (Color online) (a)First energy moment of the outgoing
Wanajo	6	electron neutrino and antineutrino as a function of time in three PNS
cooling simulations. The solid lines are the average energies of theelectron neutrinos and the dashed lines are for electron antineutrinos.

supernovae	
  can	  be	
  the	  origin	
  only	
  if	
  …	

the	
  explosion	
  is	
  not	
  due	  to	
  neutrino	
  heaIng	
  (but,	  e.g.,	
  magneto-­‐rotaIonal	
  jet;	
  Winteler+2012)	
  or	
  
our	
  knowledge	
  of	
  neutrino	
  physics	
  is	
  insuﬃcient.	
  iTHES-­‐IPMU-­‐RESCEU	

Wanajo	7	ANRV352-AA46-08

ARIr-­‐process	
  in	
  the	
  early	
  Galaxy	

15 July 2008

11:46

a

0

Relative log

–4
–6
–8
–10
– 12

log

1
0

1
0

Average abundance offsets with respect to Arlandini et al. (1999) ‘‘stellar model’’

–1
30

40

50

60
Atomic number

CS 22892-052: Sneden et al. (2003)
HD 115444: Westin et al. (2000)
BD+17°324817: Cowan et al. (2002)
CS 31082-001: Hill et al. (2002)
HD 221170: Ivans et al. (2006)
HE 1523-0901: Frebel et al. (2007)

iTHES-­‐IPMU-­‐RESCEU	

Figure 11all	
  r-­‐rich	
  GalacIc	
  halo	  stars	
  show	
  remarkable	
  agreement	
  with	  the	
  solar	  r-­‐pabern	
  v r-­‐process	
  should	
  have	  operated	
  in	
  the	
  early	
  bGalaxy;	
  
SNe	
  😃,	
  mergers	
  😢	
  ?	
  
c
v astrophysical	  models	
  should	  reproduce	  the	
  “universal”	
  solar-­‐like	  	
  r-­‐process	
  pabern	
  	
  
(for	
  Z	
  ≥	
  40;	  A	  ≥	  90)	  Sneden+2008		
  Individual stellar abundance offsets with respect to Simmerer et al. (2004)

–1

log

Annu. Rev. Astro. Astrophys. 2008.46:241-288. Downloaded from arjournals.annualreviews.org
by WIB6417 - Max-Planck-Gesellschaft on 08/04/10. For personal use only.

2

70

80

90

Wanajo	(a) Comparisons of n-capture abundances in six r-process-rich Galactic halo stars with the Solar-system r-only abundance distribution.
The abundance data of all stars except CS 22892-052 have been vertically displaced downward for display purposes. The solid light
blue lines are the scaled r-only Solar-system elemental abundance curves (Simmerer et al. 2004, Cowan et al. 2006), normalized to theEu abundance of each star. (b) Difference plot showing the individual elemental abundance offsets; abundance differences are8	constraint	
  to	
  light-­‐to-­‐heavy	
  r-­‐ra=o	Siqueira	
  Mello+…+	
  Wanajo	
  2014	
 A&A 565, A93 (2014)
2.0
1.0
0.0
-1.0
-2

-1

0

1

0

1

[Ba/Fe]
2.0
1.0
0.0
-1.0
-2

-1

VLT	
  o2.0bservaIons	
  give	  Ight	  1.5
constraint	
  for	
  light-­‐to-­‐1.0
heavy	
  0.5r-­‐abundances	
  (here	
  0.0[Sr,	  Y,	
  Zr/Ba])	
  
2
-0.5
-1.0
	
  
-1.5
v [light-­‐r/heavy-­‐r]	
  ≥-1 	
  -­‐0.3;	
   0
-2
no	
  stars	
  below	
  this	   [Ba/Fe]
2.0
constraint	
  
2

1.5
1.0

[Ba/Fe]0.5v “the	  0.0 r-­‐process”	
  must	
  
-0.5
make	
  
lighter	  r-­‐elements	
  
-1.0
with	
  
similar	
  porIon	
  
-1.5

2.0
1.0
0.0
-1.0
-2

-1

0

1

2

[Ba/Fe]iTHES-­‐IPMU-­‐RESCEU	

-2

-1

0
[Ba/Fe]Wanajo	Fig. 28. [Sr, Y, Zr/Ba] as functions of [Ba/Fe] obtained in this work
compared with results for metal-poor stars by François et al. (2007).
Symbols are the same as in Fig. 8.

9	

Fig. 29. [La/Ba] (upper panel) and [Eu/Ba] (lower pof [Ba/Fe] obtained in this work compared with resustars by François et al. (2007). Symbols are the samwww.mpa-­‐garching.mpg.de	NS	
  merger	
  scenario:	
  most	  promising?	

iTHES-­‐IPMU-­‐RESCEU	

v coalescence	
  of	
  binary	
  NSs	
  expected	
  ~	
  10	
  –	
  100	  per	  Myr	
  in	
  
the	
  Galaxy	
  (also	
  possible	
  sources	
  of	
  short	  GRB)	
  v ﬁrst	  ~	
  0.1	
  seconds	
  dynamical	
  ejecIon	
  of	
  n-­‐rich	
  maber	
  with	  Mej	  ~	  10-­‐3	
  –	
  10-­‐2	
  M¤	  v next	
  ~	
  1	
  second	  neutrino	
  or	
  viscously	
  driven	
  wind	  from	
  the	  BH	  accreIon	
  torus	
  with	
  
Mej	
  ~	
  10-­‐3	
  –	
  10-­‐2	
  M¤	
  Wanajo	10	time [d]n-induced fission is faster than radiative neutron capture at
9 K by diamonds.

Fig. 3.— Time evolution of the total radioactive heating rate

previous	
  works:	
  too	
  neutron-­‐rich	
  ?	
Mass fraction

mass fraction	e.
per unit mass, hQi, mass number hAi, and temperature hT i (allmass-averaged over the ejecta) for the 1.35–1.35M (solid lines)
final mass-integrated ejecta composition is shownand 1.2–1.5M (dotted lines) NS mergers.
.
4.
The
A
=
195
abundance
peakrelatedto2 (6pp), 2011 September 10
Goriely, Bauswein, & Janka= 126 shell closure is produced in solar distrin and found to be almost
insensitive (toallsimilar	
  input results	
  by	
  Korobkin+2011;	
  Rosswog+2013)	
Goriely+2011	
  
also	
  
eters such as the initial abundances, the expansion
100
ales, and the adopted nuclear models.
In contrast,
1.35–1.35MNS
1.35–1.35Mo NS
1.35-1.35M NS
Solar
10-1
o
eak around A = 140 originates exclusively ofrom
1.20-1.50MNSsion recycling, which takes place in the A ' 280–o
10-2gion at the time all neutrons have been captured.
nuclei are predicted to fission symmetrically as vis10-3
Fig. 4 by the A ' 140 peak corresponding to the
10-4
symmetric fragment distribution. It is emphasized
gnificant uncertainties still a↵ect the prediction of
10-5
probabilities and fragment distributions so that10-6
act strength and location of the A ' 140 fission(as well as the possible A = 165 bump observed
10-7
solar distribution) depend on the adopted nuclear
0
50
100
150
200
250
.
A
le most of the matter trajectories are subject to ay and temperature history leading to the nuclearr-­‐process	
  leading	
  ofto	  Fig.v
4.—strong	
  Final nuclear
abundance distributions
the ejecta from
1.35–1.35M (squares) and 1.2–1.5M (diamonds) NS mergers as
nd abundance distribution described above, somecycling	  53 0.60 can
0.68be0.015
0.021 0.027
0.033 0.039low
0.045-0.4 ﬁssion	  0.2atomic
0.8mass.
1.4The2.0
2.6
of
distributions
are normalized to theelements
shock-heatedat relatively
den-0.051 functions
solar
r-abundance
distribution(dottedcircles).
Typically at ⇢ > 1010 g/cm3 theYe	
Coulomb e↵ectsv severe	
  problem:	
  1.20–1.50MNS
1.20–1.50M
NSonly	  A	
  >	
  130;	
  he NSE abundance distribution towards
the ohighoescaping in neutrinos is not considered) is plotted in
Idal	
  
(or	
  2011),
weakly	  shocked)	
  ejecIon	  egion (Goriely
et al.
but atlower densities,
another	  ource	  is	
  needed	  for	
   binaFig. 3 forboth the s1.35–1.35M
and 1.2–1.5M
gh temperatures
lead
to
the
photodissociation
of
all
of	
  “pure”	  n-­‐maber	
  with	
  	
  Ye	
  <	
  0.1	
  ries. Whileandctheaverage temperature evoluthe	
  hQ(t)ilighter	  ounterpart	
  edium-mass seed nuclei into neutrons and protons.
tion di↵er only slightly between both NS-NS systems, theon recombination may occur during the decompresiTHES-­‐IPMU-­‐RESCEU	
Wanajo	11	
 velociejecta
masses Mej and mass-averaged expansion
rovided the expansion
timescale of the trajectories
ties vexp di↵er considerably. While we find for the symg enough. For a non-negligible amount of ejected
metric system vexp ⇡ 0.31c (c being the speed of light)
ial, this recombination is indeed inefficient so that

2. mergers	
  with	  GR	
  and	  ν	iTHES-­‐IPMU-­‐RESCEU	

Wanajo	12	ﬁrst	
  simula=on	
  with	  full-­‐GR	
  and	  νApproximate	
  soluIon	
  by	
  Thorne’s	  Moment	
  scheme	
  with	  a	
  closure	
  relaIon	
  } Leakage	
  +	
  Neutrino	
  heaIng	  (absorpIon	
  on	
  proton/neutron)	
  included	  }simulaIon	
  by	
  Y.	
  Sekiguchi	

Density [ log10 g/cc] 	

15

Radiation density [log10 erg/cc] 	
 30

14

29

13
12

28

11

27

10

26

9

iTHES-­‐IPMU-­‐RESCEU	

8

25

7

24
Wanajo	13	slide	
  by	
  Y.	
  Sekiguchi	

‘Robustness’	
  of	
  r-­‐process	
  in	
  	
  NS-­‐NS	
  merger	
  ?	
• Korobkin	
  et	
  al.	  2012	
  :	
  	
  	
  Ye	
  of	
  the	
  ejecta	
  depends	
  only	  weakly	
  on	
  the	
  binary	  parameters	
  so	
  that	  r-­‐process	
  in	
  the	
  NS-­‐NS	
  is	
  ‘robust’,	  – They	  adopted	
  only	  one	
  EoS	
  (Shen	
  EoS)	
  :	
  dependence	
  on	
  EoS	  is	
  not	
  explored	
  

• In	
  This	  Study	  :	
  	
  Comparison	
  between	  Steiner	
  EoS	
  and	  Shen	  EoS	  •Shen	
  EOS:	  	  	  	
  ‘S=ﬀer’	
  M NS [ M solar ]– Larger	  NS	
  radius	
  Shen	
– Mass	  ejec=on	
  is	
  driven	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  Steiner	
	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  
mainly	  by	
  Tidal	  force	
  •Steiner	
  EOS:	  	  	  	
  ‘SoRer’	
  – Smaller	
  NS	
  radius	  – Tidal	  eﬀects	
  are	
  less	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  
important	  in	
  mass	  ejec=on	
  – Stronger	
  bounce	  F ~ k EOS Δx ~ M NS ,2

RNS [km]

2−1
NS EOSE ~ k EOS (Δx) ~ M k
iTHES-­‐IPMU-­‐RESCEU	

Wanajo	14	slide	
  by	
  Y.	
  Sekiguchi	

ComposiIon	
  depends	
  on	
  EOS	
Steiner	Shen	

𝒏+​𝒆↑+ →𝒑+​𝝂 	

Higher	  T	
  :	
  more	
  	
  e+	
  	
  
	
  higher	
  Ye	
  >	
  0.25	
  region	
  :	
  	
  	
  	
  	
  	
  	
  	
  
	
  less	  neutron	
  richiTHES-­‐IPMU-­‐RESCEU	

Wanajo	Lower	  T	  :	  less	
  	
  e+	
  	
  
	
  smaller	
  Ye	
  <	
  0.25	
  :	
  	
  	
  	
  	
  	
  	
  	
  
	
  very	  neutron	
  rich
15	

10

2.0

x-y
x-z
y-z

100
-1

1.5

10-2

1.0

10-3

0.5
2

4

luminosity (1053 erg s-1)

0

4 electron ielectron anti-i
3 heavy imean energies (MeV)

10-4

25

6
8
t (ms)

1012

6
8
t (ms)

10

12

2
1
0

0

2

4

electron ielectron anti-i
20 heavy i

10
0

2

iTHES-­‐IPMU-­‐RESCEU	

4v mass	
  ejecIon	
  before	
  (40%)	
  and	  aver	
  (60%)	
  HMNS	  formaIon;	
  70%	
  ejecta	
  reside	
  near	  orbital	
  

0.0

	
  
v neutrino	
  luminosiIes	
  similar	
  between	  νe	
  and	
  anI-­‐νe	
  v neutrino	  mean	
  energies	
  similar	
  between	  νe	
  and	
  anI-­‐νe	
  

15

5

l0 (1015 cm s-3)

mass fraction (> 150 km)

neutrino	
  proper=es	
  (Steiner’s	
  EOS)	

6
8
t (ms)

10

12
Wanajo	

16	iTHES-­‐IPMU-­‐RESCEU	

Wanajo	17	nucleosynthesis	
  in	
  the	  NS	
  ejecta	
100v higher	
  and	
  wider	  range	
  of	
  	
  
Ye	
  (=	
  0.09-­‐0.45)	
  	
  
in	
  contrast	
  to	
  previous	
  cases	
  Ye	
  (	
  =	
  0.01-­‐0.05)	
  

x-y
x-z
y-z

mass fraction

10-1
10-2
10-3

10

-4

0.0

100

0.1

0.2

Ye

0.3

0.4

0.5

x-y
x-z
y-z

v higher	
  and	
  weder	  range	  of	
  entropy	  per	
  baryon	  (=	  0-­‐50)	
  in	
  contrast	
  to	
  previous	
  cases	  (=	
  0-­‐3)	
  mass fraction

10-1
10-2
10-3

10

-4

0

10

iTHES-­‐IPMU-­‐RESCEU	

20

S/kB

30

40

50
Wanajo	

18	

Ye	
  =	
  0.1	iTHES-­‐IPMU-­‐RESCEU	

Wanajo	

19	

Ye	
  =	
  0.2	iTHES-­‐IPMU-­‐RESCEU	

Wanajo	20	mass-­‐integrated	
  abundances	
10

-2

10

-3

previous	
  case	  (here	  Ye	
  =	
  0.05)	
solar r-abundance
Ye = 0.05

10solar r-abundance
this work-3

abundance

10-4

abundance

10-4
10-5

10-5

10-6

10-6

10-7

10-7

-8

10

10

this	
  work	  (for	  the	
  orbital	
  plane)	

10-2

-8

0

50

100
150
mass number

200

250

0

50

100
150
mass number

200

250

v previous	
  case:	  not	  in	
  agreement	
  with	  solar	
  r-­‐pabern	
  (e.g.,	
  for	
  	
  A	
  <	
  130)	
  à	
  also	  the	
  case	  for	
  NS-­‐NSs	
  with	
  sIﬀ	
  EOSs	  and	
  BH-­‐NSs	
  v this	
  work:	
  good	  agreement	
  with	  solar	
  r-­‐pabern	
  for	  A	  =	  90-­‐240	
  
à	
  no	
  need	
  of	
  addiIonal	
  (e.g.,	
  BH-­‐torus)	
  sources	
  for	
  light	  r-­‐elements	
  iTHES-­‐IPMU-­‐RESCEU	

Wanajo	21	3. r-­‐process	
  novae	  (or	
  goldnovae)	
  

iTHES-­‐IPMU-­‐RESCEU	

Wanajo	22	heating rate [erg g-1 s-1]r-­‐process	
  novae	
  (kilonovae)	
20 heaIng	
  rate	  from	  solar	
  r-­‐abundance	
  β-­‐decays	
1019
electro-­‐magneIc	
  transients	
  1018
Amin = 70 80 90 100 110 120 130 140 (macronova,	
  Kulkarni	
  2005;	
  
1017
1016
A = 125-135 kilonova,	
  Metzger+2010)	
  
1015
1014
	
  
1013
1012
v heaIng	
  from	  decays	
  of	
  
1011
radionuclides	
  
1010
10 9
well	
  scaled	
  as	
  dq/dt	
  ~	
  t	
  -­‐1.3	
  108
dq/dt | t-1.3
107
106
10 -6
-5
-4
-3
-2
-1
0
1
2

10

10

10

iTHES-­‐IPMU-­‐RESCEU	

10

10 10
time [d]1010

Wanajo	

10

23	

10

heatingheating

10
107
106

107
106

hea=ng	  rate	  for	  the	
  NS-­‐NS	
  ejecta	

105 -1
10

10

0

1

10

105 -1
10

2

10

100

t (d)
101

Ye = 0.09 0.14 0.19 0.24 0.29 0.34 0.39.
qsolar-r

100

10-1 -1
10

10

0

102

t (d)

.
heating rate relative to qanalytic.
heating rate relative to qanalytic

101

101

1

10

2

10

t (d)

mass-averaged `-decay fission.
qsolar-r

100

10-1 -1
10

100

101

102

t (d)

v heaIng	
  rate	  for	  the	
  mass-­‐averaged	
  abundances	
  well	  ﬁbed	  by	
  the	
  scaling	
  law	
  dq/dt	
  ~	
  t	
  -­‐1.3	
  (as	
  well	  as	
  by	
  the	
  solar	
  r-­‐pabern	
  case)	
  v but	  dependent	
  on	
  Ye;	  there	  might	
  be	
  direcIonal	
  (polar	  to	
  equatorial)	
  diﬀerences	  iTHES-­‐IPMU-­‐RESCEU	

Wanajo	24	EM	
  counterparts	
  of	
  GW	
  signals	
GW	
  signal	
  can	  be	
  spaIally	
  resolved	
  
only	
  ~	
  100	
  deg2	
  by	
  KAGRA/a.LIGO/
a.Virgo	  (from	
  2017)	
  	
  
à	
  EM	
  counterparts	
  are	
  needed	  	  v SGRBs	  events	
  should	
  be	
  restricted	
  due	  to	
  narrow	  beaming	  v r-­‐process	
  novae	  detectable	
  (by,	
  e.g.,	
  Subaru/HSC)	
  from	
  all	
  direcIons!	
  iTHES-­‐IPMU-­‐RESCEU	

Wanajo	25	already	
  found?	

LETTER

doi:10.1038/nature12505

A ‘kilonova’ associatedwith the short-duration
5″
c-ray burst GRB 130603B
N. R. Tanvir1, A. J. Levan2, A. S. Fruchter3, J. Hjorth4, R. A. Hounsell3, K. Wiersema1 & R. L. Tunnicliffe2

Short-duration c-ray bursts are intense flashes of cosmic c-rays, compact binary mergers, SGRBs may provide a bright electromagnetic siglasting less than about two seconds, whose origin is unclear1,2. The nal accompanying events detected by the next generation of gravitationalfavoured hypothesis is that they are produced by a relativistic jet wave interferometers15. Localizing electromagnetic counterparts is an
created by the merger of two compact stellar objects (specifically two essential prerequisite to obtaining direct redshift measurements and toneutron stars or a neutron star and a black hole). This is supported constraining further the astrophysics of the sources. However, theevidence supporting this progenitor hypothesis is essentially circumby indirect evidence such as the properties of their host galaxies3, N
but unambiguous confirmation of the model is still lacking. Mer- stantial: principally, many SGRBs seem to reside in host galaxies, orgers of this kind are also expected to create significant quantities of regions within their hosts, that lack ongoing star formation, which
neutron-rich radioactive species4,5, whose decay should result in a makes an origin in massive stars unlikely (in contrast to long-duration
faint transient, known as a ‘kilonova’, in the days following the c-ray bursts, which result from the core collapse of some short-lived
burst6–8. Indeed, it is speculated that this mechanism may be the massive stars16). Progress in studying SGRBs has been slow; NASA’s
E
predominant source of stable r-process elements in the Universe5,9. Swift satellite localizes only a handful per year, and they are typically
Recent calculations suggest that much of the kilonova energy should faint, with no optical afterglow or unambiguous host galaxy found inappear in the near-infrared spectral range, because of the high opti- some cases despite rapid searches with large (8-m class) telescopes.
cal opacity created by these heavy r-process elements10–13. Here we
GRB 130603B was detected by Swift’s Burst Alert Telescope onreport optical and near-infrared observations that provide strong 2013 June 3 at 15:49:14 UT17, and its duration was measured to beevidencefor such an event accompanying the short-duration c-ray
also
T90 < 0.18 6 0.02 s in the 15–350-keV band18. The burst was
iTHES-­‐IPMU-­‐RESCEU	
Wanajo	26	
burst GRB 130603B. If this, the simplest interpretation of the data, detected independently by the Konus instrument on NASA’s Windis correct, then it confirmsJ2000
that compact-object mergers are the pro- spacecraft, which found a somewhat shorter duration, T90 < 0.09 s ingenitors of short-duration c-ray bursts and the sites of significant the 18–1,160-keV band19. This places the burst unambiguously in thekilonova	
  =	
  r-­‐process	
  nova	  !	

Nature,	
  
ug.	
  2is
9	
well
e 1 | HST imaging of the location ofTanvir+2013,	
  
GRB 130603B.
TheAhost
ed and has a disturbed, late-type morphology. The position (coordinates
5 117u 049 18.299) at which the SGRB
0 5 11 h 28 min 48.16 s, dec

imaging and their difference
optical; lower row, F160W/N
with a Gaussian of width sim

r-­‐process	
  nova	  in	
  the	
  SGRB	
  aRerglow?	

Magnitude (AB)20
22

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

24

22

H

24excess	
  NIR	  indicates	
  the	  v the	  NS-­‐NS	
  ejecta	
  with	
  	
  
26
Mej	
  ~	
  0.02	  M¤	
  

26
28
30

v late-­‐Ime	
  excess	
  NIR	
  ﬂux	
   Hotokezaka et al
requires	  an	
  addiIonal	  MS1(Mej=0.07)
20
component	
  (most	  H4(M
likely	  an	
  ej=0.05)r-­‐process	
  nova)	
  APR4(Mej=0.01)rMagnitude (AB)Hotokezaka+Tanaka…+Wanajo	
  2013;	
  The Astrophysical Journal Letters, 778:L16 (5pp), 2013 November 20
NS+NS	  models	

28

0.1
1
10
Rest-frame days after GRB 130603Bv addiIonal	
  late-­‐Ime	
  red	
  30
transients	
  in	
  SGRBs	  0.1
1 should	
   10
be	
  observed	  Rest-frame days after GRB 130603B

Figure 3. Predicted light curves for NS–NS and BH–NS models. Left panel: NS–NS models. The dashed, solid, and dot-dashed curves show the H-band light curvefor the models: SLy (Q = 1.0, Mej = 0.02 M⊙ ), H4 (Q = 1.25, Mej = 4 × 10−3 M⊙ ), respectively. The total mass of the progenitor is fixed to be 2.7 M⊙ . The uppe
middle, and lower curves for each model correspond to the high-, fiducial- and low-heating models. Right panel: BH–NS models. The dashed, solid, and dot-dashe
iTHES-­‐IPMU-­‐RESCEU	
Wanajo	27	
 models are shown
curves show
the models MS1 (Mej = 0.07 M⊙ ), H4 (Mej = 0.05 M⊙ ), andAPR4 (Mej = 0.01 M⊙ ), respectively. Here only the fiducial-heating
The thin and thick lines denote the r and H-band light curves. Here we set (Q, χ ) = (3, 0.75). The observed data (filled circles), upper limits (triangles), and the lighcurves (dashed lines) of the afterglow model of GRB 130603B in r and H-band are plotted (Tanvir et al. 2013; de Ugarte Postigo et al. 2013). The observed poinin the r-band at 1 days after the GRB is consistent with the afterglow model. The key observations for an electromagnetic transient are the observed H-band data awhat	  is	
  a	
  smoking	
  gun	  of	
  the	  r-­‐process?	
ﬂux	

⽥田中，天⽂文⽉月報2014年年1⽉月
Ia	
  supernova	
r-­‐process	
  nova	hypernova	

can	
  we	
  see	  r-­‐abundances	
  in	
  
the	
  spectra?	
  	
  v almost	  featureless	
  because	  of	
  too	  many	  bound-­‐bound	
  lines	
  and	
  Doppler	  shivs	
  	
  
(v/c	
  ~	
  0.1-­‐0.3)	
  v idenIﬁcaIon	
  of	
  red,	
  featureless	
  spectral	
  wavelength	
  (Å)	
shape	
  can	
  be	
  an	
  
図 3 「r プロセス新星」の予想スペクトル．高速なunambiguous	
  evidence	  膨張により Ia 型超新星と比べて非常にのっぺ
図 4 200 Mpcの距離で起きた「of	
  an	
  r-­‐process	  り と し て い る． 極 超 新 星 と 呼 ば れ る エ ネ ル
可視光 i バンドの予想光度曲
iTHES-­‐IPMU-­‐RESCEU	
Wanajo	
ギーの高い超新星爆発ものっぺりとしたスペクトルをもつが，「r プロセス新星」は可視光に28	

星中性子星合体のシミュレ

うな高密度状態方程式を採summary	
  and	
  outlook	v NS	
  mergers:	
  very	  promising	
  site	  of	  r-­‐process	
  -­‐	
  neutrinos	
  play	  a	
  crucial	
  role	  (in	  parIcular	
  for	  a	
  sov	
  EOS)	
  v sIll	  many	  things	
  yet	  to	
  be	
  answered…	
  -­‐	
  dependence	
  on	
  mass	  raIos	
  of	
  NSs	
  and	
  EOSs;	
  how	
  about	
  BH-­‐NS?	
  -­‐	
  how	  the	
  subsequent	
  BH-­‐tori	
  contribute	
  to	
  the	  r-­‐abundances?	
  -­‐	
  can	
  mergers	
  be	
  the	  origin	
  of	
  r-­‐process	
  elements	
  in	
  the	  Galaxy?	
  iTHES-­‐IPMU-­‐RESCEU	

Wanajo	29