Mass$ejec(on$from$NS0NS$merger$ and$ … · NS0NS$merger$:$Dynamics$and$GW$waveform$ log(density$g/cc) Numerical$relavity$computaon. ... GRB$jet NS0NS$breakout GRB$jet:$$ $$$$Nakar(2007)

Mass ejec(on from NS-‐NS merger and

Kilonova/Macronova

Kenta Hotokezaka (Hebrew University of Jerusalem)

Collaborators: K. Kiuchi, T. Muranushi, Y. Sekiguchi, and M. Shibata (YITP) K. Kyutoku (UWM), H. Okawa (Waseda U), and K. Taniguchi (U. of Tokyo) M. Tanaka (NAOJ), S. Wanajo (RIKEN) T. Piran (Hebrew U. of Jerusalem)

Outline

•  Electromagne(c counterparts of Gravita(onal waves

•  Mass ejec(on at binary neutron star merger

•  A kilonova/macronova candidate associated with a short GRB 130603B

Gravita(onal-‐wave astronomy

Advanced LIGO Advanced Virgo KAGRA

GW Compact binary merger GW

Expected rate(NS-‐NS merger)

1st genera(on (Ini(al LIGO, Virgo) 2nd genera(on(Advanced LIGO, Virgo, KAGRA)

0.0002 〜　0.2 /yr 0.4 〜 400 /yr

Abadie et al (2010)

M1 =1.4Msun

M2 =1.3Msun

density

Hotokezaka, et al. (2013)

NS-‐NS merger : Dynamics and GW waveform

log(density g/cc)

Numerical rela(vity computa(on.

Waveform

Gravita(onal wave from compact binary merger

-2e-21

-1e-21

0

1e-21

2e-21

2 2.01 2.02 2.03 2.04 2.05

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time[s]

Data：Noise＋GW

-0.15

-0.1

-0.05

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0.05

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2 2.01 2.02 2.03 2.04 2.05

h

time[s]

Theore(cal Template

Compute the overlap between data and template

Matched filter analysis

Because the advanced detectors will watch 〜10^5 galaxies, many signals will be expected around the detec(on threshold.

Electromagne(c counterparts of GWs ü  Confirma(on of the detec(on of GWs from NS-‐NS around detec(on threshold, (like Neutrino burst associated with supernova 1987A) ü  Localiza(on of GW sources (GW localiza(on is not good), => Determine host galaxies.

ü  They carry different informa(on from GWs. (e.g. Mass of ejected radioac(ve nuclei)

But to discover them won’t be easy. => Theore(cal expecta(ons are needed to make observa(onal strategies.

Baryon ejec(ons drive EM counterparts

1,Dynamical ejecta: Tidal tail & shocked majer

Hypermassive NS with accre(on torus

Black hole with accre(on torus

or

2, Wind driven by viscosity, neutrino, recombina(on 3, A GRB jet may be launched at a certain (me.

Wind

GRB Jet

NS-‐NS Merger

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-‐2 0 2 4 6 8 10

log(L) [e

rg/s]

Log

Luminosity

(erg/s/Hz)

log(t) [s]

Merger remnant 　　　(radio)

Kilonova /Macronova (NIR)

Extended Emission (X)

GRB Anerglow (X)

Merger Breakout (X)

GRB Anerglow (visible)

GRB anerglow (radio)

GRB (X~γ) 52

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48

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42

log(Lν) [erg/s/Hz

]

Merger Breakout (radio)

Expected Lightcurve Refs: Nakar (2007)

Norris & Bonnell (2006) Sari, Piran, Narayan (1998)

Li & Paczynski (1998) Nakar & Piran (2012)

Kyutoku, Ioka, Shibata (2012) Kelley, Mandel, Ramirez-‐Ruiz (2012)

Luminosity

(me

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rg/s]

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(erg/s/Hz)

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GRB Anerglow (X)

Merger Breakout (X)



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]


(me

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GRB (X~γ) 52

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log(Lν) [erg/s/Hz

]


(me

Luminosity Expected Lightcurve(4π, independent of environment)

What is “kilonova/macronova”

A kilonova/macrovova was proposed by Li & Paczynski in 1998 as an observable consequence of NS-‐NS mergers.

At NS-‐NS merger ü  A frac(on of material is ejected as radioac(ve nuclei. ü  Ejecta can be bright object due to radioac(ve hea(ng. ü  Luminosity: Nova < NS-‐NS merger < Supernova.

Beta decay of radioac(ve nuclei => Keep ejecta at high T

Kilonova/Macronova and Ejecta property Luminosity

erg/s

(me

∝Mejt−1.3

Diffusion (me ∝ v−2/3M 1/3ej

Based on current understanding

〜 100 – 1000 x Nova (at the peak of a lightcurve)

〜5 days

ü  Higher ejecta mass => Brighter signal ü  Faster ejecta velocity => Brighter signal

0.01

0.1

1

10

1e-07 1e-06 1e-05 0.0001 0.001 0.01 0.1

E K/1

050er

g

M/Msun

GRNewton

GRB jet(K=30)GRB cocoon

BH-torus windHMNS wind

NS-NS breakout

Various ouylows of NS-‐NS merger

GRB jet

NS-‐NS breakout

GRB jet: Nakar (2007) GRB cocoon: Nagakura et al (2014) NS-‐NS merger breakout: Kyutoku, Ioka, & Shibata (2013) Dynamical ejecta: Hotokezaka et al. (2013) Piran et al. (2013) Bauswein et al. (2013) BH-‐torus wind: Fernandez & Metzger (2013) Just et al. (2014) HMNS – torus wind: Dessart et al. (2009) Metzger & Fernandez (2014) Perego et al. (2014)

Γ>30

Dynamical ejecta is likely dominant source of kilonova/macronova.

Numerical simula(on for dynamical ejecta

of the system for the outside of the horizon. However, thispathology could still break a numerical simulation after theformation of a black hole. To avoid this happens, weartificially set the maximum density as 1016 g=cm3 whenemploying this EOS.

Figure 2 plots the gravitational mass as a function of thecentral density and as a function of the circumferentialradius for spherical neutron stars for four EOSs. All theEOSs chosen are stiff enough that the maximum mass islarger than 1:97M!. Because the pressure in a densityregion ! & 1015 g=cm3 is relatively small (i.e., P2 issmall) for APR4 and ALF2, the radius for these EOSs isrelatively small as "11 km and 12.5 km, respectively, forthe canonical mass of neutron stars 1:3–1:4M! [38]. Bycontrast, for H4 and MS1 for which P2 is relatively large,the radius becomes a relatively large value 13.5–14.5 kmfor the canonical mass. The radius has also the correlationwith the central density !c. For APR4 and ALF2 withM#1:35M!, !c$8:9%1014 g=cm3 and !c $ 6:4%1014 g=cm3. For H4 and MS1 with M # 1:35M!, thecentral density is rather low as !c $ 5:5% 1014 g=cm3

and !c $ 4:1% 1014 g=cm3, respectively. As we show inSec. IV, the properties of the material ejected from themerger of binary neutron stars depend strongly on theradius of the neutron stars or !c.

B. Initial conditions

We employ binary neutron stars in quasiequilibria forthe initial condition of numerical simulations as in ourseries of papers [24,25]. The quasiequilibrium state iscomputed in the framework described in [39] to whichthe reader may refer. The computation of quasiequilibriumstates is performed using the spectral-method libraryLORENE [40].

Numerical simulations were performed, systematicallychoosing wide ranges of the total mass and mass ratio ofbinary neutron stars. Because the mass of each neutron star

in the observed binary systems is in a narrow range"1:2–1:45M! [38], we basically choose the neutron-starmass 1.20, 1.25, 1.30, 1.35, 1.40, 1.45, and 1:5M!. Also,the mass ratio of the observed system q :# m1=m2&' 1(where m1 and m2 are lighter and heavier masses, respec-tively, is in a narrow range"0:85–1. Thus, we choose q as0:8 ' q ' 1. Specifically, the simulations were performedfor the initial data listed in Table II.The initial data were prepared so that the binary has

about 3–4 quasicircular orbits before the onset of themerger. For four EOSs chosen, this requirement is approxi-mately satisfied with the choice of the initial angularvelocity m!0 # 0:026 for APR4 and ALF2 andm!0 # 0:025 for H4 and MS1. Here, m # m1 )m2.For the following, the model is referred to as the name‘‘EOS’’-‘‘m1’’ ‘‘m2’’; e.g., the model employing APR4,m1 # 1:2M!, and m2 # 1:5M! is referred to as modelAPR4-120150.

III. FORMULATION AND NUMERICALMETHODS

Numerical simulations were performed using anadaptive-mesh refinement (AMR) code SACRA [41] (seealso [42] for the reliability of SACRA). The formulation, thegauge conditions, and the numerical scheme are basicallythe same as those described in [41], except for the improve-ment in the treatment of the hydrodynamics code for a farregion. Thus, we here only briefly review them anddescribe the present setup of the computational domainfor the AMR algorithm and grid resolution.

A. Formulation and numerical methods

SACRA solves Einstein’s evolution equations in theBaumgarte-Shapiro-Shibata-Nakamura formalism with amoving-puncture gauge [43]. It evolves a conformal factorW :# "*1=6, the conformal three-metric ~"ij :# "*1=3"ij,the trace of the extrinsic curvature K, a conformally

0 1e+15 2e+15

M (s

olar

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!c (g/cm3) 10 15 20

M (s

olar

mas

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R (km)

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3APR4

ALF2

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H4

MS1 0

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FIG. 2 (color online). Left: The gravitational mass as a function of the central density !c for spherical neutron stars in APR4, ALF2,H4, and MS1 EOSs (the solid, dashed, dotted, and dash-dotted curves). Right: The same as the left panel but for the gravitational massas a function of the circumferential radius.

KENTA HOTOKEZAKA et al. PHYSICAL REVIEW D 87, 024001 (2013)

024001-4

We perform Numerical Rela(vity simula(ons using SACRA code Yamamoto + 2009 Solve ・Einstein equa(on ・Hydrodynamics with an Equa(on of State (4-‐different NS models) Total mass = 2.6 ~ 2.9 Msun Mass ra(o = 0.8 ~ 1

For piecewise polytropic EOSs See Read et al., (2009)

Mass ejec(on on equatorial plane

300 km x 300 km 2400 km x 2400 km

Model : 1.2Msun – 1.5Msun, APR

log(density g/cc)

300 km x 300 km 2400 km x 2400 km


log(density g/cc)

Mass ejec(on : Mej 〜 0.01Msun, v 〜 0.2c

Mass ejec(on on equatorial plane

Ejec(on Mechanism ~(dal torque~

Heavy NS

Light NS

1.  Lighter NS is elongated

2.  Outer material get angular momentum

Feature: Ejecta expand on the equatorial plane

log(density g/cc)

300 km x 150 km 2400 km x 1200 km


log(density g/cc)

Mass ejec(on on the Meridional plane (x-‐z plane)

300 km x 150 km 2400 km x 1200 km


log(density g/cc)

Mass ejec(on on the Meridional plane (x-‐z plane)

NS-‐NS Ejecta is spheroidal.

Ejec(on Mechanism ~shock hea(ng~

Model=135Msun-‐1.35Msun, APR

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H4-120150MS1-120150

FIG. 6: The central density as a function of time for models with m1 = m2 = 1.35M� (left), and m1 = 1.2M� and m2 = 1.5M�(right). Before the merger of unequal mass binaries, the central density of heavier neutron stars are plotted. �th = 1.8 isemployed for the results presented here.

FIG. 7: Snapshots of the thermal part of the specific internal energy ("th) profile in the vicinity of HMNSs on the equatorial(top) and x-z (bottom) planes for an equal-mass model APR4-135135. The rest-mass density contours are overplotted for everydecade from 1015 g/cm3.

Figures 3 – 5 indicate that there are two importantprocesses for the mass ejection. The first one is theheating by shocks formed at the onset of the mergerbetween the inner surfaces of two neutron stars. Fig-ures 7 and 8 display snapshots of the thermal part of thespecific internal energy, "th, in the vicinity of HMNSs

for APR4-135135 and APR4-120150, respectively. Thesefigures show clearly that hot materials with "th <⇠ 0.1(1.0 <⇠ 100MeV) are indeed ejected from the HMNSs,in particular, to bidirectional regions on the equatorialplane and to the polar region. This suggests that theshock heating works e�ciently to eject materials from

Specific internal energy

Spiral arm sweeps majer Mass is ejected due to the HMNS forma(on

Equatorial plane

Meridian plane

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Dependence of Ejecta mass on NS EOS

Radius of NS

Mass of Ejecta Systematics of dynamical mass ejection, nucleosynthesis, and radioactively powered electromagnetic signals 9

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|v2| [

c]

R1.35 [km]

Fig. 3.— Amount of unbound material for 1.35-1.35 M! mergers (top left) and 1.2-1.5 M! mergers (top right) for di!erent EoSscharacterized by the corresponding radius R1.35 of a nonrotating NS. Red crosses denote EoSs which include thermal e!ects consistently,while black (blue) symbols indicate zero-temperature EoSs that are supplemented by a thermal ideal-gas component with "th = 2 ("th = 1.5)(see main text). Small symbols represent EoSs which are incompatible with current NS mass measurements (Demorest et al. 2010). Circlesdisplay EoSs which lead to the prompt collapse to a black hole. The lower panels display the sum of the maxima of the coordinate velocitiesof the mass centers of the two binary components as a function of R1.35 for symmetric (bottom left) and asymmetric (bottom right)binaries.

ima of the coordinate velocities of the mass centers ofthe two asymmetric binary components. As in the sym-metric case the two stars collide with a higher impactvelocity if the initial radii of the NSs are smaller.Due to the asymmetry the dynamics of the merger pro-

ceeds di!erently from the symmetric case (see Fig. 4).Prior to the merging the less massive binary componentis deformed to a drop-like structure with the cusp point-ing to the 1.5 M! NS (top panels). After the stars beginto touch each other, the lighter companion is stretchedand a massive tidal tail forms (middle left panel). Thedeformed 1.2 M! component is wound around the moremassive companion (middle panels). Also in the case ofasymmetric mergers the majority of the ejecta originatesfrom the contact interface of the collision, i.e. from thecusp of the “tear drop” and from the equatorial surfaceof the more massive companion, where the impact ab-lates matter (see top panels). Some matter at the tipof the cusp directly fulfills the ejecta criterion (top rightpanel), while the majority obtains an additional pushby the interaction with the asymmetric, mass-shedding

central remnant and the developing spiral arms (middleright and bottom panels). A smaller amount of ejecta ofroughly 25 per cent originates from the outer end of theprimary tidal tail (particles in the lower part of the topright panel). A part of this matter becomes unbound bytidal forces (at the tip of the tidal tail in the middle leftpanel) and the other fraction by an interaction with thecentral remnant (middle left panel).Figure 5 displays the distribution of the ejecta in a

plane perpendicular to the binary orbit for the symmetricmerger (left panel) compared to the asymmetric merger(right panel) for the last timesteps shown in Fig. 2 andFig. 4, respectively. A considerable fraction of the ejectedmatter is expelled with large direction angles relative tothe orbital plane. For a timestep about 5 ms later theejecta geometry is visualized (azimuthally averaged) inFig. 6 excluding the bound matter. For both mergersthe outflows exhibit a (torus or donut-like) anisotropywith an axis ratio of about 2:3. The velocity fields alsoshow a slight dependence on the direction.

Hotokezaka + (2013)

Bauswain + (2013)

If HMNS is formed,

No massive neutron star forma(on

Similar result is obtained by MPS group.

0.0001<Mej<0.01Msun

Velocity distribu(on

0.001

0.01

0.1

0.001 0.01 0.1 1

dM/d

v

v

Most of ejecta has the velocity 0.1c 〜0.2c

Radiative Transfer Simulations for NS Merger Ejecta 9

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Fig. 8.— Expected observed ugrizJHK-band light curves (in AB magnitude) for model NSM-all and 4 realistic models. The distanceto the NS merger event is set to be 200 Mpc. K correction is taken into account with z = 0.05. Horizontal lines show typical limitingmagnitudes for wide-field telescopes (5! with 10 min exposure). For optical wavelengths (ugriz bands), “1 m”, “4 m”, and “8 m” limitsare taken or deduced from those of PTF (Law et al. 2009), CFHT/Megacam, and Subaru/HSC (Miyazaki et al. 2006), respectively. ForNIR wavelengths (JHK bands), “4 m” and “space” limits are taken or deduced from those of Vista/VIRCAM and the planned limits ofWFIRST (Green et al. 2012) and WISH (Yamada et al. 2012), respectively.


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Op(cal Near Infrared

Tanaka & KH 2013

Expected lightcurves of kilonova/macronova Masaomi-‐san’s talk in detail

〜0.01Msun

〜0.004Msun

We should follow up GW events with telescopes larger than 4m-‐size.

A Golden event: the short GRB 130603B 〜 kilonova/macronova candidate〜

ü  This could be direct evidence of compact binary merger hypothesis of short GRBs. ü  Macronovae will be promising EM counterpart of GWs. ü  A compact binary merger really produces 〜0.02Msun of r-‐process elements

Tanvir et al.,Nature,2013 Berger et al., ApJ, 2013 de Ugarte Pos(go et al, 2013

If this event is really “Kilonova/Macronova”

10−10

10−9

10−8

10−7

10−6

10−5

10−4

10−3

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Flux

den

sity

(Jy

@ 1

0 ke

V)

BAT: Black −− XRT: WT: Blue; PC: RedBAT−XRT data of GRB 130603B

0.01 0.1 1 10 100 1000 104 105 106 1070

1

2

3

K

Time since BAT trigger (s)

Short GRB130603B hjp://www.swin.ac.uk/burst_analyser/00557310/

T90 = 0.18± 0.02s

Eγ ,iso = (2.1± 0.1)×1051erg

redshin z=0.356

GRB prompt emission Swin BAT

Short GRB 130603B　

Figure 9: Light curves of GRB 130603B, indicated detections with dots and upper limits (3-�) with arrows. V -band photometry has been scaled and plotted together with the g-band. The vertical lines indicate the times whenspectra were obtained. Dotted lines indicate the light curve fits to a power law temporal decay from 0.3 to 3 daysafter the burst. We include data from the literature [21, 22]. The dashed blue line is the expected r-band lightcurve of a supernova like SN1998bw, the most common template for long GRBs after including an extinction ofAV = 0.9 magnitude. The most constraining limits indicate that any supernova contribution would be at least 100times dimmer than SN 1998bw in the r-band, once corrected of extinction (blue dashed-dotted line).

2.2 Spectral energy distribution of the afterglow and extinction

In this section we aim to fit the X-ray to optical/NIR SED using the method followed in [23, 24] to derive the

extinction in the line of sight of the GRB and determine some spectral parameters. The procedure is briefly

explained below.

The flux calibrated spectrum has been analysed after removing wavelength intervals affected by telluric lines

and strong absorption lines. We then rebinned the spectrum in bins of approximately 8 A by a sigma-clipping

algorithm. To check the flux calibration of the X-shooter spectrum, we compare the continuum with the flux

densities obtained from the extrapolation of the photometry at the time of the spectrum (mid time around 8.56 hr).

We include the X-ray spectrum from the X-ray telescope (XRT) on board Swift. We used XSELECT (v2.4) to

26

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Flux

(me

GRB anerglow(X-‐ray)

Swin XRT

Near-‐infrared excess Hubble Space telescope

Page 8 of 16

Figure 1 HST imaging of the location of SGRB 130603B. The host is well resolved

and displays a disturbed, late-type morphology. The position (coordinates RAJ2000 = 11

28 48.16, DecJ2000 = +17 04 18.2) at which the SGRB occurred (determined from

ground-based imaging) is marked as a red circle, lying slightly off a tidally distorted

spiral arm. The left-hand panel shows the host and surrounding field from the higher

resolution optical image. The next panels show in sequence the first epoch and second

epoch imaging, and difference (upper row F606W/optical and lower row F160W/nIR).

The difference images have been smoothed with a Gaussian of width similar to the psf,

to enhance any point-source emission. Although the resolution of the nIR image is

inferior to the optical, we clearly detect a transient point source, which is absent in the

optical.

u Hubble Space Telescope imaging

Op(cal

Near Infrared

9 days aner the burst 30 days

The host galaxy

Tanvir et al.,Nature,2013 Berger et al., ApJ, 2013

A macronova associated with the short GRB 130603B?

macronova candidate


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Observa(on by Hubble Space Telescope

Op(cal Near Infrared

Tanaka & KH 2013

More than 〜0.01Msun r-‐process ejecta are ejected

Hotokezaka et al ApJL (2013)

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B)

Rest-frame days after GRB 130603B

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

rH

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Rest-frame days after GRB 130603B

MS1(Mej=0.07)H4(Mej=0.05)

APR4(Mej=0.01)r

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Figure 8.3: 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 curves

for the models: SLy (Q = 1.0, Mej = 0.02M!), H4 (Q = 1.25, Mej = 4 ! 10"3M!),

respectively. The total mass of the progenitor is fixed to be 2.7M!. The upper, 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-dashed curves show the models

MS1 (Mej = 0.07M!), H4 (Mej = 0.05M!), and APR4 (Mej = 0.01M!), respectively.

Here only the fiducial-heating models are shown. The thin and thick lines denote the

r and H-band light curves. Here we set (Q, !) = (3, 0.75). The observed data (filled

circles), upper limits (triangles), and the light curves (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 point in r-band at 1 days after the GRB is consistent with the

afterglow model. The key observations for a kilonova are the observed H-band data at

7 days after the GRB, which exceed the H-band light curve of the afterglow model, and

the upper limit in H-band at 22 days after the GRB. These data suggest the existence

of a kilonova associated with GRB 130603B. This figure is taken from Hotokezaka et al.

(2013d).

124

Observed point

Expected lightvurve Mej〜0.02Msun

Expected lightcurve Mej〜0.004Msun

If dynamical ejecta are dominant contribu(on to this bump.

The observed lightcurves can be explained with Kilonova/Macronova Produced by dynamical ejecta 〜0.02Msun

GW – EM observa(on and r-‐process

1, GW observa(on when/where we should follow up. 2, EM observa(on Total mass of ejecta can be es(mated. (r-‐process element) 3, Collec(ng many events [mass/yr/galaxy] mr (r)

In order to achieve this, precise understanding of nuclear hea(ng and opacity for various type of ejecta is important.

Rebecca, Oleg, Shinya, and Masaomi’s talks

Summary Detec(on of Electromagne(c counterparts will be important. They depend on baryon ouylows. 0.0001Msun – 0.01Msun of baryons will dynamically ejected with Velocity 0.1c – 0.3c. A Kilonova/Macronova candidate associated with a short GRB 13060Bhas discovered. Es(mated dynamical ejecta 〜0.02Msun. Future We should understand possible parameter space of ejecta mass, velocity, opacity, and hea(ng rate for various type of ejecta to es(mate ejecta mass.

Documents

Mass$ejec(on$from$NS0NS$merger$ and$ … · NS0NS$merger$:$Dynamics$and$GW$waveform$ log(density$g/cc) Numerical$relavity$computaon. ... GRB$jet NS0NS$breakout GRB$jet:$$ $$$$Nakar(2007)