Dynamical mass ejection frombinary compact-star mergers
Yuichiro Sekiguchi（Toho)
with K. Kiuchi, K. Kyutoku, M. Shibata, S. Wanajo,
N. Nishimura, and K. TaniguchiToward GW astronomy

Abbott et al. 2016

Introduction

Toward GW detection and GW astronomy
GW

One of the most promising source
merger of compact-star binary with NS
Expected event rate ~ a few ×10 / yr
 Matched filter : huge parameter space

Advanced LIGO





KAGRA

Advanced Virgo

Identification of EM counterpart

Location, time, distance, etc.
Reducing parameter space (effectively S/N↑)First detection may be near the thresholdMulti-messenger confirmation

Complementary information for GW astronomy
 Gravitational waves: physics of the binary system

Mass, radius, orbit, NS EoS, BH formation

EM detection: Astrophysical environmentRedshift, host galaxy, NS EoS Dynamical mass ejection from NS-NS

Main topic

 Dynamical mass ejection from BH-NS

Discuss briefly laterlog10 [g/cm 3 ]Density Contour
in orbital plane

Tidal
InspiralMergerdeformation
Charp signal
HyperMassive NS

Gravitational Waveform

Animation by Hotokezaka
Sekiguchi et al. PRL (2011a, 2011b)
Kiuchi et al. PRL (2010); Hotokezaka et al. (2013)

Evolution of NS-NS mergers

Shibata et al. 2005,2006
Sekiguchi et al, 2011
Hotokezaka et al. 2013

Inspiral of NS binary

NS –NS merger
Dependent on
EoS, Mtot

Prompt formationof BH + Torus

Formation of hot, differentiallyrotating massive NS
Dependent on
EoS, Mtot

Delayed collapse
to BH + Torus

Rigidly rotating NS

Messengers of NS-NS mergers
GWs

Inspiral of NS binary

NS –NS merger
Dependent on
EoS, Mtot

Prompt formationof BH + Torus

Formation of hot, differentiallyrotating massive NS
Dependent on
EoS, Mtot

Delayed collapse
to BH + Torus

Rigidly rotating NS

Neutrinos

Messengers of NS-NS mergers
Inspiral of NS binary

Dynamical ejecta
R-process nucleosynthesis
External shock with ISM

NS –NS merger
Dependent on
EoS, Mtot

Prompt formationof BH + Torus

Formation of hot, differentiallyrotating massive NS

Short GRB

v-driven/MHD winds
R-process nucleosynthesis

Dependent on
EoS, Mtot

Delayed collapse
to BH + Torus
Short GRB

Rigidly rotating NS
Rotation powered
activity (like SN
remnant and pulsar)

Possible EM counterparts : Similarities to SNeSupernovaeLong GRBs












Short GRBs


Synchrotron: Ejecta-ISM interaction
Activities Powered by Pulsar

Radioactive decay of 56Ni


Merger of NS-NS, BH-NS

Prompt (γ), afterglow (X to Radio)

Supernova remnantsMerger remnants

produced in the explosive ejecta
OpticalPrompt (γ), afterglow (X to radio)

Radio Flare: Ejecta-ISM interaction
Powered by Massive NS ? (Zhang 2013)

Decay of r-process elementsProceeds in the n-rich ejecta
Opticall-IR : Macronova

Classification by spectraClassification by spectra ???



Shock breakoutMerger Shock breakout

UV ~ X. (e.g. Tominaga+ 2009)

X-ray : Kyutoku et al. (2012)

See also Metzger & Berger (2012)

SGRB and Quest for 4π emission
Jets of short GRBs may be collimated in general Jet opening angle estimated from the jet break

SGRB111020A : θj ~ 3-8° (Fong et al. 2012)



SGRB051121A : θj ~ 7° (Burrows et al. 2006)

Most of GRB Jets are expected
to be Off-Axis ⇒ very faint There will be GW-events
without SGRB counterparts

We need 4π emission events

Associated with 4π ejecta
Dynamical ejecta



neutrino-driven/MHD windsLate-time disk dissolution
 Fernandez & Metzger 2013

Rosswog @ YKIS2013

Radio flare from Ejecta-ISM interaction
External shock with inter stellar matter (ISM) : a 4π emission
 Synchrotron radiation becomes most luminous when ejecta mass =swept-up ISM mass: for typical values (Nakar & Piran 2011)1 / 3
 Eejecta   nISM   vejecta 
 


t peak ~ 4 yrs  50
3 
 10 ergs   1 cm   0.2c 
1/ 3

0.9
 Eejecta  nISM   vejecta F ~ 0.1 mJy  50
3  
 10 ergs  1 cm   0.2c 

2.8

5 / 32   obs 
D

 

200
Mpc
1.4
GHz





ISM density may be much smaller : according to recent SGRB obs.
nISM ~ 0.01-0.1 cm-3 for SGRB 111020A (Fong et al. 2012)
 nIMS ~ 0.0001-1 cm-3 for SGRB 111117A (Margutti et al. 2012)




0.75

Radio flare may be less bright and shine in a very late time :Not very suited as EM counterparts of GWs

Rotation powered activities ?


If a stable massive NS is survived, additional EM emissions powered by
NS-rotation may be expected (Metzger et al. 2011; Zhang 2013; Gao et al. 2013)Compared to normal pulsars, rapid rotation (P~ms), strong B-fields (B~1015 G)

However, such additional emissions may be not very frequent :Nuclear theory : hard to make such a very stiff EoS with Mmax > 2.4Msolar SGRB : if central engine of SGRB is BH + Disk, frequent formation of the massive
NS may lead to too much mergers (only low mass NS merger ?) ~1/3 of SGBRs may have late-time activity
Gao et al. 2013which could be originated in the massive SN

Most of them are short duration < O(100s)Collapse to a BH ?
shorter than the spin down timescale > 1000s

Macronova
Merger ejecta will be very neutron rich: rapid neutron capture (r-process)proceeds (Lattimer & Schramm 1974) : n + (Z,N) ⇒ (Z,N+1)
 Competition with the β-decay :(Z,N+1) ⇒ (Z+1,N) + e + νeThe r-process is very sensitive to how much neutrons are there, that is, tothe electron fraction Ye ( = Yp = 1 – Yn) : we need michrophysics !

Then, EM transients powered by radioactivity of the r-process elements
are expected (Li & Paczynski 1998; Kasen et al. 2013; Tanaka & Hotokezaka 2013) v t peak ~ 10 days 

0
.
3
c



1 / 2

1/ 2

M




0
.
01
M
solar






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
2
10
cm
/
g



1/ 2

1/ 2

1/ 2

M
 f  v  Lpeak ~ 10 41 erg/s  6   10  0.3c   0.01Msolar 
1/ 4

 f   v 
eff
Tpeak
~ 2 103 K  6   10   0.3c 

1 / 8






2
 10 cm / g 

M


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
 0.01Msolar 1 / 8

1 / 2

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


2
 10 cm / g 

3 / 8

Importance of Ye in the r-processElectron fraction (Ye) is the key parameter : Ye ~ 0.2-0.25 is critical thresholdYe < 0.2-0.25 : strong r-process ⇒ nuclei with A>130
Ye > 0.2-0.25 : weak r-process ⇒ nuclei with A< 130

Different decay heat and opacity for them (Smaller κ for smaller A: Grossman et al. 2013

Kasen et al. 2015)


Neutrino-matter interactionYe can be changed
Two reactions which increase YePositron capture : n  e   p  eImportant for higher temperature∵ there are more positrons Neutrino capture : n   e  p  eStrong Yedependence

Copious neutrinos are emitted
 NS matter is neutron rich
 Not considered in the previous
studies (need neutrino transfer)Korobkin et al. 2012

Key observations for r-process : Universality
solar pattern = event by event r-process pattern




Solar

Sneden et al. (2008)

Abundance patterncomparison :

r-rich low metallicity starsSolar neighborhood

Low metallicity suggests
Such stars experience a
few r-process events
 Such stars preserve theoriginal pattern of the
r-process events(chemical fossil)Key observations for r-process : Universality
solar pattern = event by event r-process pattern
The solar and chemical
fossil r-process element
patterns agree well suggests that


Solar

r-process event synthesize
heavy elements with apattern similar to solar
pattern (Univsersality) Low metallicity !!
 Universality should beachieved before chemicalenrichment

Sneden et al. (2008)

Should not rely on many
eventsSingle event has to satisfy
the universality

Dynamical mass ejection from NS-NSWith ‘Universality’ point of view :NS-NS merger ejecta: too neutron-rich ?
Goriely et al. 2011; Bauswein et al. 2013
Approx. GR SPH sim. without weak interactions
 No way to change Ye => ejecta remains n-rich (initial low Ye)
 See also post-process calculation of weak interactions


Korobkin et al. 2012; Rosswog et al. 2013
Newtonian SPH sim. with neutrino
 tidal mass ejection (explained in the next slide) of ‘pure’ neutron star matter

Ejecta is very n-rich with Ye < 0.1 ??

Mass ejection from BNS merger (1) :Tidal torque + centrifugal force
Hotokezaka et al. (2013)

Less massive NS is
tidally deformedAngular momentum
transfer by spiral arm
and swing-by



A part of matter isejected along theorbital planereflects low Ye of
cold NS (β-eq. at T~0),
no shock heating,rapid expansion(fast T drop), no time
to change Ye by weakinteractions
Density contour[ log (g/cm3) ]With ‘Universality’ point of view :NS-NS merger ejecta: too neutron-rich ?
Korobkin et al. 2012; Rosswog et al. 2013; see also Goriely et al. 2011
tidal mass ejection of ‘pure’ neutron star matter (very n-rich) with Ye < 0.1

Ye is that of T=0, β-equilibrium

strong r-process with fission recycling only 2nd (A~130; N=82) and 3rd (A~195;
N=126) peaks are produced (few nuclei in A=90-120)
the resulting abundance pattern does not satisfy universality in A=90-120

T=0, β-eq.

Korobkin et al. (2012) MNRAS 426 1940

1st peak

2nd

3rd

Goriely et al. (2011) ApJL 738 32

How to satisfy the universalityElectron fraction (Ye) is a key parameter : Ye ~ 0.2 is critical threshold

Ye < 0.2 : strong r-process ⇒ nuclei with A>130 (the pattern is robust)
Ye > 0.2 : weak r-process ⇒ nuclei with A< 130 (for larger Ye, nuclei with smaller A)

We need ejecta
with higher Ye
Korobkin et al. 2012

How to satisfy the universality
Introduce new ejecta components

Neutrino driven winds from the remnant systemlate time disk/torus disintegrationDessart et al. (2009); Grossman et al. (2014); Perego et al. (2014); Just et al. (2015)

Fernandez & Metzger (2013)

It is not clear whether it is possible to satisfy the universality robustly

 Take into account effects of both GR and weak interaction in

the dynamical ejecta (this talk)What will change if you include GR and microphysics (1) :Stronger shock in GR
van Riper (1988) ApJ 326 235

SN shock velocity @ 300 km(1000km/s)

realistic ‘stiffness’e.g., Kolehamainen et al. (1985)
NPA 439 535

General relativisitic：
Stronger shock wave formation

Newtonian gravity :Weaker shock and its independence of EOS

Shock is stronger in GR
Shock velocity dose notdepend on EOS in
Newtonian gravity (EOS
is already too stiff )Incompressibility (‘Stiffness’ of EOS) K(sym) (MeV)Massejection from BNS merger
(2):
0
0
0
10
20
30
40
0
10
20
Shock drivent (ms) components
t (ms)

30

40

G . 6: T he central density as a function of tim e for m odels with m 1 = m 2 = 1.35M ( left), and m 1 = 1.2M and m 2 = 1.5M
 Shocks occur due to oscillations of massive NS and collisions of spiral armsght). Before the m erger of unequal m ass binaries, the central density of heavier neutron stars are plotted. Γ t h = 1.8 isployed for the results presented here.

Isotropic mass ejection, higher temperature (weak interactions set in)Specific internalenergyx-yx-z

Hotokezaka et al. (2013)

What will change if you include GR and microphysics (1) :Stronger shock in GR
Newtonian simulation by S. Rosswog et al.

Almost no isotropic component(shock-driven) in Newtonian
simulation
Only the tidal component

Full GR simulation by Y. Sekiguchi et al.

What will change if you include GR and microphysics (2) :Ye can change via weak interaction
Driven by tidal interactions

Consists of cold NS matter in
β-equilibrium ⇒ low Ye and TDriven by shocks

Consists of shock heated matter
higher temperature =>Weak interaction can change Ye

x-z

Previous studies and our study

Korobkin et al. 2012 : Newtonian SPH simulations with neutrinos

Bauswein et al. 2013: Relativistic SPH simulations with many EOS but without neutronosThis Study : Full GR, approximate gray radiation hydrodynamics simulation withmultiple EOS and neutrinos (brief summary of code is in appendix of lecture note) Einstein’s equations: Puncture-BSSN/Z4c formalism
BNS merger simulations with multiple EOS,
 GR radiation-hydrodynamics (neutrino heating can be approximately treated)different
mass ratio, and switch-on and –off
 Advection terms : Truncated Moment scheme (Shibata et al. 2011)
 EOS : any tabulated
EOS with 3D smooth connection
to Timmes
EOS
neutrino
interactions
becomeaccessible
 gray or multi-energy but advection in energy is not included
thankstoand
XC30
CfCA and K-computer Fully covariant
relativisticin
M-1 closureSource terms : two options
Implicit treatment : Bruenn’s prescription
Explicit treatment : trapped/streaming ν’s










e-captures: thermal unblocking/weak magnetism; NSE rate
Iso-energy scattering : recoil, Coulomb, finite size
e±annihilation, plasmon decay, bremsstrahlung
diffusion rate (Rosswog & Liebendoerfer 2004)
two (beta- and non-beta) EOS method

Lepton conservation equations

Neutrino energy density

Adopted EOS & (expected) Mass ejection mechanism‘Stiffer EOS’










⇔ RNS : larger
TM1, TMA
Tidal-driven dominant
Ejecta consist of low T & Ye
NS matter‘Intermediate EOS’

© M. Hempel

DD2
TM1

IUFSU

DD2

TMA

‘Softer EOS’







⇔ RNS : smaller
SFHo, IUFSU
Tidal-driven less dominant
Shock-driven dominantYe can change via weakprocesses

SFHo

See also, Bauswein et al. (2013); Just et al. (2014)

Entropy per baryon : DD2
relatively stiff, tidal component dominatedYe : DD2
relatively stiff, tidal component dominatedEntropy per baryon : SFHo
relatively soft, multiple shock components

Ye : SFHo
relatively soft, multiple shock components

34Sekiguchi et al PRD (2015)

Soft(SFHo) vs. Stiff(TM1): Ejecta temperature
Soft (SFHo): temperature of unbound ejecta is higher (as 1MeV) due tothe shock heating, and produce copious positrons Stiff (TM1): temperature is much lower
Soft (SFHo: smaller RNS)

1000kmStiff (TM1: larger RNS)k BT ~ a few MeV  me c 2  0.511 MeV

𝒏 + 𝒆+ → 𝒑 + 𝝂Higher T : more e+
Shock heating
more positron capture

Lower T : less e+
Mass ejection mainly
driven by tidal effects

Sekiguchi et al PRD (2015)

Soft(SFHo) vs. Stiff(TM1): Ejecta Ye = 1- Yn
Soft (SFHo): In the shocked regions, Ye >> 0.2 by weak processes
 Stiff (TM1): Ye is low as < 0.2 (only strong r-process expected)
Soft (SFHo: smaller RNS)

Stiff (TM1: larger RNS)𝒏 + 𝒆+ → 𝒑 + 𝝂Higher T : more e+higher Ye > 0.25 region :less neutron rich

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

Wanajo, Sekiguchi et al. ApJL (2014)

Achievement of the universality(soft EOS (SFHo), equal mass (1.35-1.35))
010mass fraction

10

x-y
x-z
y-z

-1

10

-2

10

-3

10

-4

0.0

0.1

0.2

10

20

Ye

0.3

0.4

0.5

30

40

50

0

10

-1

mass fraction

10

x-y
x-z
y-z

10-2
10-3
10-4



0

S/kB

The Ye-distribution histogram has a broad, flat structure (Wanajo, Sekiguchi, et al. (2014). )

Mixture of all Ye gives a good agreement with the solar abundance !
Robustness of Universality (dependence on binary parameters)

Sekiguchi et al PRD (2015)

EOS dependence : 1.35-1.35 NS-NS
Softer EOS

Softer EOS

Ejecta Ye
Dynamical ejecta mass
Softer EOS
 Mej is larger for softer EOS

Consistent with piecewise-polytrope studies
 Only SFHo will give Mej ~ 0.01 Msun a value required by the total amountof r-process elements and flux of the‘macronova’ event (GRB 130603B)

Ye distribution

Unequal mass NS-NS system: SFHo1.25-1.45


Orbital plane : Tidal effects play a role, ejecta is neutron rich
Meridian plane : shock + neutrinos play roles, ejecta less neutron rich

Dependence on binary parameterfor soft EOS (SFHo)

30ms after merger

Sekiguchi et al PRD (2015); Prego et al. (2014); Just et al. (2014); Goriely et al. (2015); Martin et al. (2015)

Importance of neutrino absorption in
dynamical mass ejection

Ejecta Ye
Dynamical ejecta mass

Amount of ejecta mass can beincreased order of 10-3 Msun



Average Ye can change 0.02~0.03
depending on EOS : effect isstronger for stiffer EOS where
HMNS survive in a longer time

Ye distribution

Evolution of BH-NS

Shibata & Taniguchi (2008)
Kyutoku et al. (2010), (2011)

rtidal < rISCO
plungertidal > rISCO

ISCO: innermost
stable circular orbit

rISCO  rISCO ( M BH , aBH )BH spin dependence
larger aBH ⇒ smaller rISCO

Tidal radius :tidal force  self gravity of NS
tidal disruption
Imre Bartos, GECo, Columbia University(Bartos et al. 2013, in prep.)rtidal  M NS  M BH  M BH 

2/3

 M NS  RNS 

1

Compactness of NS⇒ NS structure (EOS)Lattimer & Prakash (2001) ApJ 550 426

Symmetry energy and NS radiusNS radius is sensitive tosymmetry energy (@ρs)
empirical correlation forradius and pressure near
the saturation density
RNS (nS , M NS ) 
1/ 4

 P ( nS ) C (nS , M NS )
-3 
 MeV fm For pure neutron matter,P @ saturation depends on
symmetry energy

P(nS , x  0)  L  SV nS

S potential : fixedSV  S K  S potential

S K : fixedSymmetry energy of adopted EOS

SHFo : L=47 MeV
DD2 : L=55 MeV
TM1 : L=111 MeV

R 4  nS SV  nS L / 3

BH-NS merger (DD2 EOS: density)
MBH=5.5Msun, MNS=1.35Msun, aBH=0.75
 aBH-NS merger (DD2 EOS: Ye)
MBH=5.5Msun, MNS=1.35Msun, aBH=0.75

Properties of ejecta：EOS dependenceR  14.5 km, L  111 MeV

R  13 km, L  55 MeV

R  12 km, L  47 MeVMain mass ejection mechanism: tidal disruption of NS by BH ⇒ejecta is very neutron richFor lager RNS (for stiff EOS or EOS with larger symmetry energy), ejecta mass is larger(promising as EM counterpart to GW)

Neutrino heating effects less dominant
Too neutron rich to satisfy the Universality

Need some additional ejecta components. (e.g., MHD driven winds: Kiuchi, YS, et al. 2015)

Not only ejecta mass but also ejecta Ye
reflects the symmetry energyR  14.5 km, L  111 MeV

R  12 km, L  47 MeV

NS is tidally disrupted :ejecta 〈Ye〉 reflects NS matterYe in beta equilibrium