The r-process in neutron star mergers
Shinya Wanajo (NAOJ)
Y. Sekiguchi (YITP), N. Nishimura (Keele Univ.),
K. Kiuchi, K. Kyutoku, M. Shibata (YITP)
Supernovae and Gamma-Ray Bursts 2013
October 14-November 15, 2013, Kyoto University

contents	

1. supernovae vs. mergers
2. mergers with GR and ν	

3. r-process novae (kilonovae)

SNGRB2013	

Wanajo	2	1. supernovae vs. mergers

SNGRB2013	

Wanajo	3	origin of gold (r-process elements) is still unknown…

SNGRB2013	

Wanajo	4	popular r-process scenarios	

core-collapse supernovae(since Burbidge+1957;
Cameron 1957)v n-rich ejecta nearby proto-NSv not promising according torecent studies
SNGRB2013	

Wanajo	neutron-star mergers(since Lattimer+1974;
Symbalisty+1982)v n-rich ejecta from coalescing
NS-NS or BH-NSv few nucleosynthesis studies
5	ANRV352-AA46-08

ARI

r-process in the early Galaxy	

15 July 2008

11:46

a

0

Relative log

–4
–6
–8
–10
– 12

log

1
0

all r-rich Galactic halostars show remarkableagreement with the solarr-patternv r-process should haveoperated in the earlybGalaxy;SNe 😃, mergers 😣
cv astrophysical models
should reproduce the“universal” solar-liker-process pattern(for Z ≥ 40; A ≥ 90)
Sneden+2008	

Individual stellar abundance offsets with respect to Simmerer et al. (2004)

–1
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.

20

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)

70

80

90

HD 221170: Ivans et al. (2006)
HE 1523-0901: Frebel et al. (2007)

SNGRB2013	

Figure 11

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 are6	recipe to cook gold	
requisite physical conditionsv high entropy; S (∝ T 3/ρ)to slow seed productionv short expansion timescale; τ(e-folding time of T)
to freeze seed productionv low electron fraction; Ye(protons per nucleon)
to leave free neutrons

SNGRB2013	

Wanajo	7	how to make the 3rd peak and beyond	

physical condition for making A ≥ 200 (Hoﬀman+1997)v high entropy case (S > 100 kB nuc-1)
f200 = (S / 230 kB nuc-1)3 / [(Ye / 0.4)3 (τ / 20 ms)] ≥ 1v low entropy case (S < 100 kB nuc-1);
Ye < 0.2
with any S, τ	

SNGRB2013	

Wanajo	8	8.8 M¤ self-consistently

exploding ONeMg core supernova

	simulation by Bernhard MüllerSNGRB2013	

Wanajo	9	SNGRB2013	

Wanajo	10	SN ejecta: not so neutron-rich…	

PHYSICAL REVIEW C 86, 065803 (2012)

IUM MODIFICATION OF THE CHARGED-CURRENT . . .

ννeo illustrate the effect of the correct inclusion of meaneffects in charged-current interaction rates, as well as thertance of thev
nuclear
PNS coolingfor symmetry
Ye < 0.5energy,(i.e.,fiven-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
SNGRB2013	
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

no	
  correc3on	

(a)
0.6
0.55
Ye,NDW

Δ = M − M ≈ 1.29 MeV

12

8

ε
Ye ~ %1+( ,$ Lν e εν e + 2Δ 'III. PROTO-NEUTRON STAR EVOLUTIONnpνe~ 6 MeV	ν16<εν> (MeV)

While collisional broadening tends to increase both νe andoss sections,vRPAdecrease thebyνe cross secYecorrelations
is determinedand 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 theseequilibriumisra, we set aside
effects due to value
RPA correlations
andsional broadening, and calculate the neutrino interactions−1 asincluding the mean field energy shifts calculated#&Lν e ν e − 2Δ
ribed in Ref. [1].

no

n	
o
3
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	electron neutrino and antineutrino as a function of time in three11	
PNS
cooling simulations. The solid lines are the average energies of theelectron neutrinos and the dashed lines are for electron antineutrinos.

SN scenario: not very promising?	
Wanajo, Janka, Müller 2011; ECSN	

-2

10

10

solar r-abundanceunchanged
Ye,min changed

-3

10

-3

10

M = 1.2 1.4 1.6 1.8 2.0 2.2 2.4 Msun
MNS	

solar r-abundance

10-4
10-5

-5

10

10-6

10-6
10-7

10-7

-8

10-8

10
10

Wanajo 2013; semi-analytic model	

abundance

abundance

10-4

-2

Ye,min = 0.40

0.30

0.20

0.15

-9

10

-9

0

50

100
150
mass number

200

250

50

100

150
mass number

200

250

SNGRB2013	

to solar r-abundance102
wind (S ~ 100-200)
early dynamical ejecta (S ~ 10-30) neutrino-driven
M = 1.2 1.4 1.6 1.8 2.0 2.2 2.4 Msunv 2D self-consistent SN models v recent wind simulations found
1
10Y
found Ye > 0.4 (Wanajo+2011);
e > 0.4 (Roberts+2012); only
weak
r-process (Wanajo 2013)
only up to A = 90
0
Wanajo	1012	to make a long story short, …	

core-collapse supernovae appear not to be the r-process site
as far as the mass ejection is due to neutrino heating.

SNGRB2013	

Wanajo	13	www.mpa-garching.mpg.de	

NS merger scenario: most promising?	

SNGRB2013	v coalescence of binary NSs
expected ~ 10 – 100 per Myr in
the Galaxy (also possible sourcesof short GRB)v dynamical ejection of n-richmatter with Mej = 10-4 – 10-2 M¤(e.g., Bauswein+2012;
Hotokezaka+2013)v neutrino- (or viscous, MHD)
wind from the BH accretion torus
with Mej ~ 10-3 M¤(e.g., Surman+2008; Metzger+2010; Wanajo & Janka 2012)
Wanajo	14	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 ratee.
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 (alsoto all similar
input results by Korobkin+2011; Rosswog+2013)	
Goriely+2011
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 nuclear
r-processleading to
Fig.v
4.—strong
Final nuclear
abundance distributionsof 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 severe1.20–1.50M
problem:
only A > 130;
1.20–1.50M
NSNS
he NSE abundance distribution towards
the ohighoin neutrinos is not considered) is plotted in
tidal(or2011),
weaklyshocked)ejectionescapinganotheregion (Goriely
et al.
but atlower densities,
is needed
for binaFig. 3 for both thesource
1.35–1.35M
and 1.2–1.5M
gh temperatures
lead
to
the
photodissociationofallof “pure” n-matter with Ye < 0.1 ries. Whileand theaverage temperature
evoluA <hQ(t)i
130 (notuniversal)??
edium-mass seed nuclei into neutrons and protons.
tion di↵er only slightly between both NS-NS systems, theon recombination may occur during the decompresSNGRB2013	
Wanajo	15	
 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

Mass fraction

mass fraction	

previous works: too neutron-rich ?	

there are still shortcoings …	

general relativisty was not considered or only approximately included,
and neutrino transport was not taken into account at all.

SNGRB2013	

Wanajo	16	2. mergers with GR and ν	SNGRB2013	

Wanajo	17	our new (and first) work is…	

study or r-process nucleosynthesis based on a
fully general-relativistic hydro simulation withapproximative neutrino transport.

SNGRB2013	

Wanajo	18	slide	
  by	
  Y.	
  Sekiguchi	Summary	
  of	
  implementa3on	• Einstein’s equations: Puncture-BSSN/Z4c formalism• GR radiation-hydrodynamics (Sekiguchi + 2013)	
  – Advec&on	
  terms	
  :	
  Truncated	
  Moment	
  scheme	
  (based	
  on	
  Shibata	
  et	
  al.	  2011)	
  	
  • Fully	
  covariant	  and	
  rela3vis3c	  	
  • gray	  or	
  mul3-­‐energy	
  but	
  advec3on	
  in	
  energy	
  is	
  not	  included	  • M-­‐1	
  closure	
  • EOS	  :	
  any	
  tabulated	
  EOS	
  with	  3D	
  smooth	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  	
  
connec3on	
  to	
  Timmes	
  EOS	  

– Source	
  terms	  :	
  	
  two	
  op3ons	  • Implicit	
  treatment	
  :	
  Bruenn’s	
  prescrip3on	  	
  • Explicit	
  treatment	
  :	
  	
  trapped	
  /streaming	  	
  ν’s	
  
–
–
–
–
–

e-­‐captures:	
  thermal	
  unblocking/weak	
  magne&sm;	
  NSE	
  rate	  	
  Iso-­‐energy	
  scaBering	
  :	
  recoil,	  Coulomb,	
  ﬁnite	  size	  e±annihila&on,	
  plasmon	
  decay,	
  bremsstrahlung	  	
  diﬀusion	
  rate	  (Rosswog	
  &	
  Liebendoerfer	
  2004)	
  two	
  (beta-­‐	
  and	
  non-­‐beta)	
  EOS	
  method	  – Lepton	
  conserva&on	
  equa&ons	
  Wanajo	

SNGRB2013	

19	

slide	
  by	
  Y.	
  Sekiguchi	Leakage	
  +	
  neutrino	
  hea3ng	

Approximate	
  solu3on	
  by	
  Thorne’s	  Moment	
  scheme	
  with	  a	
  closure	
  rela3on	
  } Leakage	
  +	
  Neutrino	
  hea3ng	  (absorp3on	
  on	
  proton/neutron)	
  included	  }

Density [ log10 g/cc] 	

15

Radiation density [log10 erg/cc] 	
 30

14

29

13
12

28

11

27

10

26

9

SNGRB2013	

8

25

7

24
Wanajo	20	slide	
  by	
  Y.	
  Sekiguchi	Temperature	
  depends	
  strongly	
  on	
  EOS	
Steiner	Shen	

𝒏+​𝒆↑+​𝒆↑→+𝒑	+​𝝂 	Higher	  T	
  :	
  more	
  	
  e+	
  	
  
	
  Shock	
  hea&ng	
  	  	  more	  positron	
  capture	  	
  SNGRB2013	

Lower	  T	  :	  less	
  	
  e+	
  	
  	
  
	
  Mass	  ejec&on	
  mainly	
  	
  	
  	
  	
  
	
  driven	
  by	
  &dal	  eﬀects	
Wanajo	21	slide	
  by	
  Y.	
  Sekiguchi	Neutrino	
  emission	
  depends	
  on	
  EOS	
Steiner	Shen	

Larger NS radius &Stiffer EOS : smaller
emissivity due tolower T

Smaller NS radius &Softer EOS : larger
emissivity due tohigher T
SNGRB2013	

Wanajo	22	slide	
  by	
  Y.	
  Sekiguchi	Neutrino	
  emission	
  depends	
  on	
  EOS	
Steiner	Shen	

𝒏+​𝒆↑+ →𝒑+​𝝂 	

Higher	  T	
  :	
  more	
  	
  e+	
  	
  
	
  lager	  ​𝝂   emissivity

SNGRB2013	

lower	  T	  :	  less	
  	
  e+	
  	
  
	
  smaller	
  ​𝝂   emissivity

Wanajo	23	slide	
  by	
  Y.	
  Sekiguchi	

Composi3on	
  depends	
  on	
  EOS	
Steiner	Shen	

𝒏+​𝒆↑+ →𝒑+​𝝂 	

Higher	  T	
  :	
  more	
  	
  e+	
  	
  
	
  higher	
  Ye	
  >	
  0.25	
  region	
  :	
  	
  	
  	
  	
  	
  	
  	
  
	
  less	  neutron	
  richSNGRB2013	

Wanajo	Lower	  T	  :	  less	
  	
  e+	
  	
  
	
  smaller	
  Ye	
  <	
  0.25	
  :	
  	
  	
  	
  	
  	
  	
  	
  
	
  very	  neutron	
  rich24	

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:	  	  	  	
  ‘So`er’	
  – 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
SNGRB2013	

Wanajo	25	our result: not so neutron-rich ?	
Ye	

3D, full-GR, ν-transport
simulation of
a 1.3 + 1.3 M¤ merger
with the “soft” EOS of
Steiner+2013v wide range of
Ye ~ 0.1-0.4 becauseof e-, e+ captures andneutrino eﬀects
v S ~ 10-30 kB/nuc;Ye < 0.2 is needed for
r-processing
simulation by Y. Sekiguchi	SNGRB2013	

Wanajo	26	108

neutrinos increase Ye	
107 -5
10angle-averaged neutrino luminosities	

10-1

100

101

0.5
0.4

with iwithout i

2.5

0.37
0.3

2

Ye

Luminosity [1053 erg/s]

10-2
t (s)

Ye evolution for a selected trajectory	

ele-nuele-anu
heavy-nu

3

10-3

5GK

3.5

10-4

1.5

0.2

0.19

1

0.1

0.5
0
0

5

10
15
Time [ms]2025v neutrino fluxes several times
greater than those of SNe

SNGRB2013	

0.0 -5
10

10-4

10-3

10-2
t (s)

10-1

100

v neutrinos substantially
increase Ye

Wanajo	27	101Ye = 0.1	

SNGRB2013	

Wanajo	28	Ye = 0.2	

SNGRB2013	

Wanajo	29	nucleosynthesis in the NS ejecta	
100

10-2
10-3

abundance

10-4
10-5

10-2

10-6

10-3

10-7

-4

0.0

0.1

0.2

0.3

0.4

10

0.5

Ye

v post-process nucleosynthesis
calculations for 2D(equatorial) 36 trajectories
SNGRB2013	

-8

0

50

100
150
mass number

200

250

10-2

solar r-abundancev variation of r-processes
mass-averaged
10
depending on Ye
-3

10-4

abundancemass fraction

10-1

10

Ye = 0.09 0.14 0.19 0.24 0.29 0.34 0.39 0.44

Wanajo	

10-5
10-6

30	

10
10

-8

0
50
100
150
200
250
mass-integrated
abundances	mass number
10-2

solar r-abundancemass-averaged

10-3
abundance

10-4
10-5
10-6
10-7

10

-8

0

50

100
150
mass number

200

250

v good agreement with the full range of the solar r-patternv strong support for NS mergers being the origin of r-elements
SNGRB2013	

Wanajo	31	our work concludes …	

neutrinos cannot be ignored for r-process calculations.

SNGRB2013	

Wanajo	32	3. r-process novae (or goldnovae)

SNGRB2013	

Wanajo	33	kilonova: smoking gun of an r-process?	

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
SNGRB2013	
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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,e 1 | HST imaging of the location ofTanvir+2013,
GRB 130603B.
TheAug.
host29	
is well
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

you say kilonova, I say r-process nova	
Hotokezaka+2013; NS+NS models	

Magnitude (AB)20
22

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

22

APR4(Mej=0.01)rHv heating rates from radioactive
24
decays of r-process products(as26well as the opacities andejecta masses) determine the28brightness

24
26
28
30

r-process novae (or goldnovae):Hotokezaka et alimportant EM counterpartsofMS1(Mej=0.07)
20
H4(Mej=0.05)
GW sources
Magnitude (AB)The Astrophysical Journal Letters, 778:L16 (5pp), 2013 November 20

0.1
1
10
Rest-frame days after GRB 130603B

à30talk
0.1 by M. Tanaka
110Rest-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
SNGRB2013	
Wanajo	35	
 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1012
heating rate (erg g-1 s-1)

1020
1019
Ye = 0.09 0.14 0.19 0.24 0.29 0.34 0.39 0.44181017
| t-1.3
1016
1015
1014
1013
1012
1011
1010
10 9
108
107
106
10 -6
10 10-5 10-4 10-3 10-2 10-1 100 101 102
time (d)1011

Ye = 0.09 0.14 0.19 0.24 0.29 0.34 0.39 0.44
| t-1.3

1010
109
108
107 -1
10

100

101

102

time (d)

1020
1012
19
-1.3 (Metzget+2010)10v
approximately follows ~|t t-1.3
18
1017
11
10
10v
within a factor 2 range for Ye < 0.3;
1016
1015
ejecta of Ye > 0.3 do not significantly
1010 contribute to heating
1014
1013
12
9
10SNGRB2013	
Wanajo	
 10
111010
10 9 mass-averaged
mass-averaged
108 mass-averaged (`-decay)
108
eating rate (erg g-1 s-1)

eating rate (erg g-1 s-1)

heating rate (erg g-1 s-1)

heating rate for each Ye	

| t-1.336	

heating r

10
108

heating rate for all the ejecta	

1020
1019
| t-1.3
18
1017
1016
1015
1014
1013
1012
1011
1010
10 9 mass-averaged
108 mass-averaged (`-decay)
107 solar r-abundance (`-decay)
106
10 -6
10 10-5 10-4 10-3 10-2 10-1 100 101 102
time (d)107 -1
10

100

101

1012

| t-1.3

1011
1010
109
108
107 -1
10

mass-averaged
mass-averaged (`-decay)solar r-abundance (`-decay)
100

101

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time (d)v fission contribution is unimportant in our case;β-decay heating assuming the solar r-pattern does a good job(used in Hotokezaka+2013; Tanaka+2013)v factor 2 deficiency because of half the ejecta having Ye > 0.3
SNGRB2013	

102

time (d)heating rate (erg g-1 s-1)

heating r
heating rate (erg g-1 s-1)

1011
1010
10 9
108
107
106
10 -6
10 10-5 10-4 10-3 10-2 10-1 100 101 102
time (d)37	summary and outlook	v NS mergers: very promising site of r-process- neutrinos play a crucial role (in particular for a soft EOS)v still many things yet to be answerd…
- dependence on mass ratios of NS/NS; how about BH-NS?
- detailed spatial information of radioactive heating?
- can mergers be the origin of r-process elements in the Galaxy?
SNGRB2013	

Wanajo	38