RIKEN-IPMU-RESCEU Joint meeting, 140707

r-process enrichment by neutronstar mergers
Based on Tsujimoto & TS 2014

Kilonova
The Astrophysical Journal Letters, 774:L23 (4pp), 2013 September 10

✤

✤

✤

Berger+ 13, Tanvir+ 13!
NIR detection!
no detection of optical band

Kilonova (M =0.1, v =0.2c)
ej

H

−20

Kilonova (M =0.01, v =0.2c)
ej
ej
−2.6
)
Afterglow (F ∝ t

i
r
g

−18
−16
−14

ej

ν

SN2006aj−12

HST

Afterglow of GRB130603B atday 9!

−22

Magellan/Gemini✤

NIR counterpart of neutronstar merger (short GRB)!
Absolute Magnitude

✤

−10
−8
−1

10

0

10
Rest−Frame Time (d)110Figure 2. Absolute magnitude versus rest-frame time based on our ground-based
observations from Magellan (Section 2), on Gemini data (Cucchiara et al. 2013),
and on our HST photometry (Section 2; blue: F606W; red: F160W). Also shown
is an afterglow model with a single power law decline of Fν ∝ t −2.6 , required

Kilonova✤

NIR counterpart of neutronstar merger (short GRB)!

✤

Afterglow of GRB130603B atday 9!
✤

Berger+ 13, Tanvir+ 13!

✤

NIR detection!

✤

no detection of optical band

Kilonova

EARCH LETTER
Time since GRB 130603B (d)
1
21

10
X-ray

10–11

F606W
22

F160W

AB magnitude

23

10–12

24
25
10–13

26

X-ray flux (erg s–1 cm–2)

eld. The redshifts of the afterglow21 and the host galaxy22 wereound to be z 5 0.356.
other proposed signature of the merger of two neutron stars or a
✤ NIR counterpart of neutron
on star anda black hole is the production of a kilonova (somealso termeda ‘macronova’or(shortan ‘r-processsupernova’) due tostarmergerGRB)!
cay of radioactive species produced and initially ejected during
erger process—in other words, an event similar to a faint, short6–8. Detailed calculations suggest that the spectra ofsupernova✤ Afterglow of GRB130603B atkilonova sources will be determined by the heavy r-process ions
10–13
are
d in the day
neutron-richmaterial.
Althoughthesemodels9!
r from being fully realistic, a robust conclusion is that the optical
ill be greatly diminished by line blanketing in the rapidly expanejecta, withthe radiation emerging instead in the near-infrared✤ Berger+ 13, Tanvir+ 13!
and being produced over a longer timescale than would otherbe the case. This makes previous limits on early optical kilonova
ion unsurprising23. Specifically, the NIR light curves are expected
e a broad✤peak,risingdetection!
after a few days and lasting a week or more
NIR
rest frame. The relatively modest redshift and intensive study of130603B made it a prime candidate for searching for such a kilonova.
imaged oflocationof the burst with
the NASA/ESA
Hubble✤ thenodetectionofopticalband
Telescope (HST) at two epochs, the first ,9 d after the bursth 1) and the second ,30 d after the burst (epoch 2). On each occaa single orbit integration was obtained in both the optical F606W
0.6 mm) and the NIR F160W filter (1.6 mm) (full details of the imag-

27
28

10–14

29
104

105
106
Time since GRB 130603B (s)Figure 2 | Optical, NIR and X-ray light curves of GRB 130603B. Left axis

GRB130603B
The Astrophysical Journal Letters, 774:L23 (4pp), 2013 September 10✤

Bershort duration GRB!
✤

T90~0.18±0.02 s in 14-350
keV!

✤

z=0.356

Figure 1. Magellan and HST observations of the afterglow and host galaxy of GRB 130603B. Top: Magellan/IMACS and LDSS3 r-band
36.4 days, respectively, with the location of the afterglow marked. Middle: HST/ACS/F606W images. Bottom: HST/WFC3/F160W im
circle marks the location of GRB 130603B, with a radius of 10 times the rms of the astrometric match between the Magellan and HSTfading source coincident with the afterglow position is clearly visible in the WFC3/F160W image, with no corresponding counterpart in th(A color version of this figure is available in the online journal.)Neutron star merger v.s. Kilonova✤

Mass ejection from neutron star merger!
✤

Bauswein+ 2013!
✤

SPH simulation (GR)!
✤1.35+1.35 M⦿ neutron star coalescences (1.2+1.5)!

✤

various EOSs!

✤

ejects matter enriched by r-process elements!
✤

0.001-0.01 M⦿ at speeds of ~0.2cBauswein, Goriely, & Janka
14.5

30

14.5

14

14
20

13.5

13.5

13

11.5
11

−10

12.5
12

0

11.5
11

−10

The AstrophysicalJournal, 773:78 (21pp), 2013 August 10
10.5

10.5
10

−20

10

−20

9.5
−20

−10

0
x [km]

10

20

30

9.5
−30
−30

9

12.1235 ms
14.5

30

−20

−10

0
x [km]

10

20

30

30

20

13
10

12

0

11.5

y [km]

12.5

11

−10

12.5

11

10

−20

−10

0
x [km]

10

20

30

−30
−30

9

−20

−10

0
x [km]

10

20

30

9

40

14

40

14 0.32

30

13.5

30

13.5

13

20

13 0.31
10
12.5

−20
−30

12.5

11.5
11
10.5
10

10
0
−10
−20
−30

9.5

12

13
R1.35 [km]

14

15

13
[km]

14

15

0.345
0.34
0.335
0.33
0.325
0.32
11

12

13
14
R
[km]

15

16

1.35

0.315
11

12

R

1.35

Figure 3. Amount
12of unbound material for 1.35–1.35 M⊙ mergers (top left) and 1.2–1.5 M⊙ mergers (top right) for different EOSs characterized by theradius R1.35 of a nonrotating NS. Red crosses denote EOSs which include thermal effects consistently, while black (blue) symbols indicate zero-temp
11.5 by a thermal ideal-gas component with Γth = 2 (Γth = 1.5) (see main text). Small symbols represent EOSs which are incompatibthat are supplemented
NS mass measurements(Demorest et al. 2010; Antoniadis et al. 2013). Circles display EOSs which lead to the prompt collapse to a black hole. The
11
display the sum of the maxima of the coordinate velocities of the mass centers of the two binary components as a function of R1.35 for symmetric (boasymmetric (bottom
10.5 right) binaries.
(A color version of this figure is available in the online journal.)
10

9.5 2013), whereas the relativistic calculations of
2013; Rosswog
this study andin Oechslin et al. (2007) and Goriely et al.
−50
−50
9
9
−50
0
−50
0
50
(2011) 50
yield only a few 10−3 M⊙ of unbound material for thex [km]
x [km]
15.16713.8024 ms
1.35–1.35
M⊙ msbinary with the same EOS. Comparing the resultsdensityof ourstudy with the likewise relativistic calculations in
2. Merger and mass ejection dynamics of the 1.35–1.35 M⊙ binary with the DD2 EOS, visualized by the color-coded conserved rest-mass(logarithmicallyin g cm−3 ) in the equatorial plane. The dots mark SPH particles which represent ultimately gravitationally unbound matter. Their
positions are projectionsof
Hotokezaka
et al. (2013)shows very good agreement for all
ee-dimensional locations anywhere in the merging stars onto the orbital plane. Black and white indicate the origin from the one or the other NS. For every tenthfour EOSs used in Hotokezaka et al. (2013). For example, fore the coordinate velocity is indicated by an arrow with a length proportional to the absolute value of the velocity (the speed of light corresponds to a line length
−40

0.01

0.005
11

16

0.34

50

−10

13.0056 ms

12

15

9.5

14.5

0

13
14
R1.35 [km]

0.35

10

−20

50

10

12

0.36

13.4824 ms0.33
14.5

20

11

10.5

9.5
−30
−30

0
10

11.5

−10

0.015

0.002

12

0

10.5
−20

0.006

13.5 0.004

13
10

0.01

9

14

13.5

0.02

12.6867 ms0.008
14.5

14
20

0.012

Mejecta [Msun]

−30
−30

Bauswein, Gori

Mejecta [Msun]

12y [km]

12.5

0

1.35+1.35

1310|v1|+|v2| [c]

10

2

20

|v1|+|v | [c]30y [km]

y [km]

y [km]

y [km]

strophysical Journal, 773:78 (21pp), 2013 August 10

−40

−3

EOS dependence of the ejecta production is still
particular with a consistent description of thermal ef
given EOS the radius R1.35 of a nonrotating NS wi
is a characteristic quantity specifying the compactn
Therefore, we use R1.35 to describe the influence odensity EOS on the amount of NS merger ejecta.
The upper left panel of Figure 3 displays t

The Astrophysical Journal, 773:78 (21pp), 2013 August 10

Figure 8. Nuclear abundance pattern for the 1.35–1.35 M⊙ mergers with the
NL3 (blue), DD2 (red), and SFHO (green) EOSs compared to the solar r-process
abundance distribution (black).
(A color version of this figure is available in the online journal.)

Bauswein, Goriely, & Janka

Figure 9. Nuclear abundance pattern for the 1.2–1.5 M⊙ mergers with the NL3(blue), DD2 (red), and SFHO (green) EOSs compared to the solar r-process
abundance distribution (black).
(A color version of this figure is available in the online journal.)

Most of matter composed of r-process elements

Results of Nucleosynthesis

a few high-density EOSs (two EOSs were used in Goriely et al.
2011). Since the NS EOS affects sensitively the dynamics of NS
mergers and thus the properties of the ejecta (amount, expansion
velocity, electron fraction, temperature), we explore here theinfluence of the NS EOS on the r-process nucleosynthesis in asystematic way.
For a selected, representative set of EOSs we extract the日付thermodynamical
histories of fluid elements which get gravitationally unbound. For these trajectories nuclear network calculations were performed as in Goriely et al. (2011), where details

nuclei), which are still subject to large uncertainties (Goriely
et al. 2009b). Hence, the deviations from the solar abundance
pattern between A ≈ 130 and A ≈ 170 are not unexpected,
while the third r-process peak around A = 195 is a consequenceof the closed neutron shell at N = 126, which is robustly
predicted by theoretical models. Very similar results wereobtained for NS merger models performed with the LS220 andShen EOSs in Goriely et al. (2011).
In Figure 9 the normalized abundance patterns are shown forasymmetric 1.2–1.5 M⊙ mergers employing the same represen-Light from r-process
dominated ejecta
The Astrophysical Journal, 775:18 (9pp), 2013 September 20✤

Barnes & Kasen 2013!
✤

opacity of r-process elements >>opacity of Fe!
✤

but little is known!

✤

Assume r-process elements havethe same opacity as Nd.!

✤

Lanthanide(Z=58-72)!

✤

Line blanketing !
✤Faint in optical (BVR)!

✤

bright in NIR (H)Light from r-process
dominated ejecta

ophysical Journal, 775:18 (9pp), 2013 September 20

Barnes & Kasen

The Astrophysical Journal, 775:18 (9pp), 2013 September 20✤

Barnes & Kasen 2013!
✤

opacity of r-process elements >>opacity of Fe!
✤

but little is known!

✤

Assume r-process elements havethe same opacity as Nd.!

✤

Lanthanide(Z=58-72)!

✤

Line blanketing !
✤Faint in optical (BVR)!

Bolometric light curves for βchar = 0.1 and a range of ejecta masses,
using line data from two different Autostructure models of Nd,a somewhat different
energy level structure. Solid curves correspond✤ bright
in(opt3),
NIR while(H) dashed curves show
eferred Autostructure modelm an alternative model (opt2). The moderate differences can be taken
estimate of the error resulting from uncertainties in the Nd structure

Tanaka &Hotokezaka 2013

strophysical Journal, 775:113 (16pp), 2013 October 1

Tanaka & Hotokezaka

42

10

Similar simulations!
NSM-all✤

NSM-dynamical
NSM-windopacity for all r-process
NSM-Feelements with Z>31!

1041

✤

✤

40
10
✤

κ~10 cm2/g!

Monte Carlo (Lucy 2005)!

Similar1 results

10Days after the merger

e 4. Bolometric light curves for simple models with different elemental
ances: NSM-all (31 ! Z ! 92), NSM-dynamical (55 ! Z ! 92),

Absolute magnitude✤-18

UBV RI J HK Bol NSM-all

-16
-14
-12
-10
-8
-6

2

4 6 8 10 12 14 16 18 20
Days after the merger

Figure 5. Multi-color light curves of the model NSM-all (in Vega magnitudes).
The light curves in redder bands are brighter and decay more slowly.
(A color version of this figure is available in the online journal.)

Short summary✤

detection of H-band emission by HST is Kilonova !

✤sGRB is a result of neutron star merger!

✤✤

ejects r-process rich matter!

✤

Mass=0.001-0.01 M⦿!

✤speeds~>0.1c!

Motivation to revisit NSM as a major production site of r-processelements

Rarity of NSM in faint dSphs

[Ba/H]r-process element
in dSphs
no evolution or sudden
increase of Eu in small mass
dSphs !
✤✤✤✤

s-process enrichment?
-3
Hercules-5

Ursa Major II

-4

-3.5

-3

-2.5

-2

-1.5

[Fe/H]

1

faint & massive dSphs

0.5

very few NS merger events!

Eu/Fe tells us the merger
rate relative to the SN rate~
one merger per 1,000-2,000
SN if 0.01 M⦿ of Eu supplied
per event-2

ComBer[X/Y]=log(X/Y)-log(X/Y)⦿!

Eu/H increase with increasing
Fe/H in massive dSphs!

Ultra
Faint
ultra faint
dSphs
dSphs

-4

Sgr

0

[Eu/H]

✤

-1

-0.5
-1

Leo I

Fornax

Draco

-1.5
-2
-2.5
-2.5

Sculptor

Carina-2-1.5[Fe/H]

-1

-0.5

0

r-process elements in extremely
metal-poor stars(EMP stars) in
the Milky Way✤

Large dispersion in [Eu/Fe]!

✤Supernova supplies ~0.1 M⦿Fe!
✤

SNR sweeps and mixes ISM
of mass ~105 M⦿: [Fe/H]~-3!

✤

Some EMP stars are thesecond generation

From SAGA (Suda+ 08)

r-process elements in extremely
metal-poor stars(EMP stars) in
the Milky Way✤

Large dispersion in [Eu/Fe]!

✤Supernova supplies ~0.1 M⦿Fe!
✤

SNR sweeps and mixes ISM
of mass ~105 M⦿: [Fe/H]~-3!

✤

Some EMP stars are thesecond generation

From SAGA (Suda+ 08)

kilonova remnant✤

Energy ~1050 erg!

✤

ejected mass of Ba~10-4-10-2 M⦿!

✤Swept up mass ~104 M⦿ (Fluid approximation is applied)!

✤If a star forms from the swept up gas: [Ba/H]~0-2!

✤[Fe/H] determined by other SNe!

✤If [Fe/H]=-3, then [Ba/Fe]~3-5 too high !!?

Transferproperties
RJ\✤!"#" $%&'()*& +) (," - ./0," 123)%" (24The speed of ejecta = 0.1-0.3c(γ=1.010-1.099) !
✤Stopping length ~ 400/n kpc @0.2c
Fluid approximation is bad!

✤

Energy per
nucleon=muc2(γ-1)=10-100 MeV/A!

✤

If 0.3c, then spallation occurs
before traveling through the
stopping length!

✤

If 0.1c, then below the thresholdenergyW(<9 \9 @#',. !"#$$%$&!'(#)$ 8#" 5"#'#)>2",)(2+ !#..($(#) ,$ ,82)!'(#) #8 &)&"<-9 @4& &=5&"(+&)',. 5#()'$ ,"& ',b&) 8"#+
AB\*BKC9 @4& $#.(1 .()& "&5"&$&)'$ #2" "&$2.'$9

Effects of magnetic fields✤✤

If B~1 $G!
✤

rg~2×1015 cm<<size of a galaxy!

✤

Structure of B decides the dissipation of kinetic energy of r-processelements!

✤

ordered B fields does not stop r-process elements!

✤

turbulent B may stop r-process elements through ionizations!

r-process elements propagate along tangled B fields and pervade theentire proto-galactic cloudr-process elements in EMP stars✤

in ~7/n Myr after a single event of NS merger!
✤

r-process elements are uniformly distributed!
✤✤

EMP stars form in gas swept up by individual SNe!
✤✤concentration of r-process elements is determined by the cloudmass!

Fe abundance is determined by the SN!

High(Low) [r-process/Fe]→low (high) cloud mass

Chemical
evolution model✤

Fe is distributed in a swept up shellwith the mass of 105 M⦿!
2.5

✤

✤

r-process elements diffuse over theentire proto-galactic cloud.!
NSM rate: dSphs!
Chemical evolution in clouds withdifferent masses!
Blue: 109 M⦿!
✤ Red: 107-2×107 M !
⦿✤ Green: 2×105, 2×106 M
⦿✤[Eu/Fe]

✤

2
1.5
1
0.5
0
-0.5
-1
-1.5

-4

-3

-2

[Fe/H]-10

r-process
Fig. 9. Comparison of the [Ba/H], [La/H] and [Eu/H] values with the stelsynthesis fitting
errors. Errors on [Eu/H] are the spectrum synthesis fitting
elements
inof the two Eu lines summed in quadrature with the fitting error.
globular cluster
s
T

eff

(K)

log g

−1.4

✤

−1.8

−1.8

−2.0

−2.0

variation in [Eu/H] and−2.2[Ba/H]!
−2.4

✤

−2.6
no significant
variation in−2.8[Fe/H]!
−3.0

✤

−1.6

−2.4 −2.2 −2.0 −1.8 −1.6 −1.4 −1.2 −1.0
[Eu/H]

massive GC!

[La/H]

M15!

a)[Ba/H]

✤

−1.6

−1.4

b)

−2.2
−2.4
−2.6
−2.8
−3.0

−2.4 −2.2 −2.0 −1.8 −1.6 −1.4 −1.2 −1.0
[Eu/H]

Fig. 10. a) [Ba/H] and b) [La/H] compared with [Eu/H] for the W13✤ core collapsed!
(black dot) sample compared with S97 (blue plus), S00a (red square),
O06 (blue triangle), and Sk11 (RGB: red cross, RHB: pink asterix) sam✤ hosts
plessome
wheredouble
available.
pulsars

t
i

o
F
d
p
w
h
–
t
h
a
s

Fig. 12. Histograms exploring the three key samples of M 15 giants in [Ba/H]: a) all 63 W13 stars (blue) and 24 bright W13 stars (green); b) all
57 DO10W13 stars (blue) and 21 bright DO10W13 stars (green); and c) the W13 sample combined with the further 12 S97W13 stars and the further
19 DO10W13 stars with oﬀsets applied.

Table 10. Mean and standard deviations of [Fe/H], [Ba/H], [Eu/H], and [La/H] derived (where possible) for the three samples of W13, DO10W13
and S97W13 , for each whole sample and separated into the respective low (I) and high (II) Ba modes.

W13
DO10W13
S97W13
W13 Mode I
W13 Mode II

N
63
57
18

[Fe/H]NBa[Ba/H]
NEu[Eu/H]NLa[La/H][Ba/Eu][La/Eu]–2.34 ± 0.06 63 –2.20 ± 0.26 20 –1.70 ± 0.13 13 –2.04 ± 0.16 –0.50 ± 0.29 –0.24 ± 0.21
–2.33 ± 0.06 57 –2.47 ± 0.31 −
−
−
−
−
−
–2.33 ± 0.03 18 –2.36 ± 0.28 18 –1.78 ± 0.17 −−–0.58 ± 0.33−

[Ba/H] < –2.20 –2.36 ± 0.06 30 –2.41 ± 0.16 7 –1.80 ± 0.08
[Ba/H] ≥ –2.20 –2.33 ± 0.05 33 –2.00 ± 0.16 13 –1.65 ± 0.13

DO10W13 Mode I [Ba/H] < –2.47 –2.33 ± 0.05 29 –2.74 ± 0.18
DO10W13 Mode II [Ba/H] ≥ –2.47 –2.34 ± 0.07 28 –2.22 ± 0.16
S97W13 Mode I
S97W13 Mode II

−
−

−
−

[Ba/H] < –2.36 –2.34 ± 0.01 10 –2.56 ± 0.12 10 –1.92 ± 0.15
[Ba/H] ≥ –2.36 –2.33 ± 0.05 8 –2.11 ± 0.19 8 –1.76 ± 0.18

58

–2.19 ± 0.13 –0.61 ± 0.18 –0.39 ± 0.15
–1.95 ± 0.11 –0.35 ± 0.21 –0.30 ± 0.17

−
−

−
−

−
−

−
−

−
−

−
−

–0.64 ± 0.19
–0.35 ± 0.26−−

Notes. The number of stars (N) use to calculate the element abundance.

compared to the mean of Mode II for W13. The spread in Eu values within each mode is of similar magnitude and not significantly smaller than that of the spread in the Ba values. While
La abundances are only available for W13 the spread on thewhole sample and for each mode are again similar in magnitude
to the corresponding spread in Eu.
Certainly for each sample and each element, there is a clear
separation in mean values between the modes. Even for Eu andLa the separation between the means of Mode I and Mode II are
outsidethe 1σ limit.
日付[Ba/Eu] nor [La/Eu] with [Eu/H]. This is not unexpected, as if
all the contribution from the r-process has been removed thenthe s-process contribution would not trend with increasing rprocess abundances. This is borne out by the distribution lying
closer to the estimated solar r-process only yield abundance (indicated as red dashed lines, Simmerer et al. 2004) than to thesolar abundance (red dotted lines), which is a combination ofboth r- and s-process contributions. However the distribution
does not lie exactly on the estimated solar r-process only yield
implying that there is some contribution of the heavy elementsfrom some other process. This is replicated in Fig. 13a for the
S97W13 sample.

Summary of M15 observations

Stellar surface polluted with NSM
ejecta?
✤Red-giants with r-process excess!
✤Possible scenario !
✤First star formation event!

✤all supernovae went off!
✤

all the gas component removed!

✤less massive stars evolved to AGB and supplied gas!

✤

NSM enriched the gas with r-process elements!

✤

still less massive dwarfs have accreted the r-process rich gas✤some dwarfs become extremely
r-process rich!
the dwarfs evolved to red-giants
to be observed!

✤

dN/d[Eu/H]✤

15

The distribution with respect tor-process elements becomes
bimodal.

10

5

0
-2.4

-2.2

-2

-1.8
[Eu/H]

-1.6

-1.4

-1.2

✤

some dwarfs become extremely
r-process rich!
the dwarfs evolved to red-giants
to be observed!

✤

dN/d[Eu/H]✤

15

The distribution with respect tor-process elements becomes
bimodal.

10

5

0
-2.4

-2.2

-2

-1.8

[Eu/H]

-1.6

-1.4

-1.2
-1.2

✤

some dwarfs become extremely
r-process rich!

✤

The distribution with respect tor-process elements becomes
bimodal.

lim

to be observed!

10

Pr
e

the dwarfs evolved to red-giants
dN/d[Eu/H]

✤

in
ar

y

15

5

0
-2.4

-2.2

-2

-1.8

[Eu/H]

-1.6

-1.4

-1.2
-1.2

Future prospects✤

gravitational wave detections from NS mergers!
✤

sensitive to events closer than short duration GRBs!

✤

follow up observations (γ-ray, X-ray, optical, NIR)!
✤

direct measurement of spectral sign of r-process elements!
✤✤

estimate of mass from light curve analyses!
✤✤

abundance and velocities!

Origin of r-process elements!

EOS of dense matter