Radioactively Powered Emission
from Neutron Star Mergers(kilonova, macronova, r-process nova, goldnova, ...)

Masaomi Tanaka (NAOJ)- MT & Hotokezaka 2013, ApJ, 775, 113- MT, Hotokezaka, Kyutoku, Wanajo, Kiuchi, Sekiguchi, Shibata
2013, ApJ, in press (arXiv:1310.2774)- Hotokezaka, Kyutoku, MT, Kiuchi, Sekiguchi, Shibata, Wanajo
2013, ApJ, 778, L16

EM emission fromEM counterparts of=NS-NS merger
GW sources

strophysical Journal, 746:48 (15pp), 2012 February 10Metzger & Berger

specific stellar populations). Because merger counterparts
Jet ISM Shock (Afterglow)
redicted to be faint, obtaining a spectroscopic redshift
Optical (hours days)
Radio (weeks years)
allenging (cf. Rowlinson et al. 2010), in which caseroscopy of the host galaxy is the most promising means
Ejecta ISM ShockRadio (years)
taining the event redshift.
s important to distinguish two general strategies for conobs
ng EM and GW events. One approach is to search for aGRB(t ~ 0.1 1 s)signal following an EM trigger, either in real time or at
t-processing stage (e.g., Finn et al. 1999; Mohanty et al.
Kilonova). This is particularly promising for counterparts predicted
Optical (t ~ 1 day)cur in temporal coincidence with the GW chirp, such asjMerger Ejecta-duration gamma-ray bursts (SGRBs). Unfortunately, mostTidal Tail & Disk Wind
promising counterparts (none of which have yet beenv ~ 0.1 0.3 c
endently identified) occur hours to months after coalesBH.6 Thus, the predicted arrival time of the GW signal willn uncertain, in which case the additional sensitivity gained
this information is significantly reduced. For instance, if
me of merger is known only to within an uncertainty ofurs (weeks), as we will show is the case for optical (radio)
erparts, then the number of trial GW templates that must
Metzger & Berger 2012,arched is larger by a factor ∼104 –106 than if the merger
Figure 1. Summary of potential electromagnetic
counterpartsof NS–NS/ApJ,746,
48
is known to within seconds, as in the case of SGRBs.
NS–BH mergers discussed in this paper, as a function of the observer angle
second approach, which is the primary focus of this paper,θobs . Following the merger a centrifugally supported disk (blue) remains aroundthe S.
centralcompact object and(usuallyB.
a BH).
Rapid accretion
lasting !1 s
M follow-upof GW triggers.
potential advantage
in thisN. Gehrels’s,
A.AFruchter’s,
S. Wanajo’s,
Rosswog’s,
Metzger’s
talks
powers a collimated relativistic jet, which produces a short-duration gamma
is that counterpart searches are restricted to the nearby• On-axis short GRB• Off-axis radio/opticalafterglow• Radioactive emission(r-process nuclei)

2 deg Kiso Supernova Survey (KISS)
Kiso 1m Schmidt telescope~ 200 galaxies / 4deg2(< 200 Mpc)

GW alert error box
e.g. 10 deg x 10 deg~ 5000 galaxies(< 200 Mpc)

2 deg

179

186

193

201

208

215

223

230

237

244

251

Expected emission
Li & Paczynski 98
Kulkarni 05
Metzger+10
A. Bauswein’s and B. Metzger’s talks

Timescale

Luminosity

tpL

1 day

M
0.01M

1042 erg s 1

energydepositionenergydeposition

1/2

M
0.01Mv
0.2c
1/2

1/2

1/2

0.1 cm2 g 1v
0.2c

1/2

1/2

0.1 cm2 g 1

Opacity of Fe~ 20 mag at 200 Mpc(1m-class telescopes)

10

Opacity of Lanthanoid Ce

Kasen, BCeκ = 10 cm2 g-1
κ = 1 cm2 g-1
2-1
κ = 0.1 cm g

FeFig. 9.— Comparison of the expansion opacity for pure cer
High opacity (xlines)
100)
Lowerluminosity(xcalculations1/10) ass
and=>the VALD
linelist(green lines). The
wavelength binning= 0.01 . Left: Line expansion opacity v
Kasenetal.
2013
the Autostructure results are in reasonably good agreement witBarnes& Kasenironejecta (brownline). 2013
Right: Planck mean opacity as a functRadiative transfer in NS merger ejecta
with all the r-process elements

Numerical relativity

3D, time-dependent,
multi-frequency
radiative transfer

Hotokezaka et al. 2013

MT & Hotokezaka 2013

1. Opacity of r-process-dominated ejecta?
2. Characteristic feature of NS merger?
3. Dependence on EOSs and progenitors?

Solar mass fraction

106

Number of lines

Radiative Transfer Simulations for NS Meger Ejecta

104

105
104
103
105

TotalIII
III
IV

103
102
101
100

0510152025303540 45 50 55
Atomic number Z

6065707580859095ig. 1.— (Upper) Solar abundance r-process abundance raio in mass fraction (Simmerer et al. 2004). The abundance is normalized wGe) = 106 . (Lower) The number of bound-bound transition data for different elements. Different colors show different ionization statm neutral (I) to triply ionized ion (IV). The atomic data at Z ≤ 30 are taken from Kurucz & Bell (1995) while the data at Z ≥ 31
mpiled using the VALD database (Piskunov et al. 1995; Ryabchikova et al. 1997; Kupka et al. 1999, 2000). It is shown that there ista for triply ionized ions (IV, purple) at Z ≥ 31.

~500,000 transitions (up to Fe)+100,000 transitions (r-process)cket is assigned as tactive = t0,decay z −5 , so that it reoduces the energy release following t−1.2 . Here t0,decay
the beginning of the radioactive energy release. In thisper, we set t0,decay = 10−4 days, which is sufficiently
rlier than the initial time of the simulations. Whene computation reaches to t > tactive , UVOIR packets
e created (instead of γ-ray packet for the case of SNe,neglecting the γ-ray transfer, see Section 3). Similar

consider the electron scattering, and free-free, boun
free, and bound-bound transitions. The wavelengt
dependent opacity is evaluated in each cell after the teperature estimate in each time step. The bound-bou
transition is the dominant source of opacity both fType Ia SNe and NS mergers.
Electron scattering:
By solving the Saha equtions, the number density of free electrons (ne ) is coputed in each cell. The absorption coefficient of electr

Evaluate bound-bound opacity(in each time step)As in10has a hopacity54). T
NSM-all(cm2 g-171) ha2 )-1
κ=0.1 cm gopacityκ = 0.1
κ=1.0 cm2 g-1
that, hκ = 12 -1
the opa
κ=10 cm g
κ = 10Figu
41
NSM-a10is red b(2) a h
Full simulation
opticalmatica
reddersistent
MT & Hotokezaka 2013
Kasen
40
10Since
110ments,ulation
Days after the merger
includeFig. 2.— Bolometric light curve of model NSM-all (black, multicuss thHigher
opacity
byfactorof100
frequency simulations). It is compared with the light curves for thespectramodel
but with the grayapproximationbyof UVOIRtransferof 10Faintersame
thanpreviously
expectedafactorafter th(κ =0.1, 1, and 10 cm2 g−1 for blue, purple, and red lines, respec(consistentKasen+13, Barnes
& are
Kasen
all thetively).
The resultwithof multi-frequency
transfermost13)
similar to-1

UVOIR Luminosity (erg s )42

Log flux (Fh) + constant

UB
VRI
65Type Ia

JH
1.5 days
5.0 days10.0 days

432
1

5000

10000
15000
20000
Wavelength (A) MT & Hotokezaka 2013

- Very red SED (peak at NIR)
- Extremely broad-line (feature-less) spectra(Identification of r-process elements seems difficult)Log flux (Fh) + constant

UBVRI
6J

54

SN 1998bw

H
1.5 days
5.0 days10.0 days

32
1

5000

10000
15000
20000
Wavelength (A) MT & Hotokezaka 2013

- Very red SED (peak at NIR)
- Extremely broad-line (feature-less) spectra(Identification of r-process elements seems difficult)

42

10

1.2 + 1.5
1.3 + 1.4

toy model
NSM-all
APR4 (soft)
SoftEOSH4 (stiff)(ug bandsRadiusof1.35 deepMsunobser
NS-1

UVOIR Luminosity (erg s )lengths (iR = far
11.1below
km tForkmlarge-2
R=
13.6
10 MsunStiff EOS
be smaller
41
10
Hawaii Teltelescope/
3.6 deg2 an
the limitinbright mod40
10
at the firstservations
faint mod10-3 Msun
below theMT & Hotokezaka 2013posure).
39
10To cover1
10class telesc
aru/Hyper
Days after the merger
and Large
Fig. 7.— Bolometric light curves for realistic models (Table 1).
2008;
LSS
SofterNSmass
ratio
The
expected EOS/Higheremission of models witha soft
EOS APR4(red) iswide fieldbrighter than that with a stiff EOS H4 (blue). For the soft EOS
We show
Brighteremission
APR4, the light curve
does not dependon the mass ratio, while for
optical wa
a stiff EOS H4, a higher mass ratio (1.2M! + 1.5M! , solid line)(Hotokezaka+13;
Bauswein+13,M. Shibata’s,
A.emission
Bauswein’s
results in a large
ejecta mass, andthus, brighter
than atalks)can detect
lower mass ratio (1.3M+ 1.4M , dashed line).

20020Observing strategy after GW detection
20ndMpc

At 200 Mpc
fsot

22

10 min5 sigmaS

24i band
optical
200 Mpc
i band1m1m

EO

234m4mff
EO

25SObserved magnitude

21istnd

51015
Days after the merger

5days
26

8m8m

27
20

051015
Days after the merger

20 Best with i band (0.8 um)
NIR (~21 AB mag in NIR)
J band20

2NS-NS
Hotokezaka+13

8

HR diagram
brighter-20-18

SN Ia
SN IIPSN IbcBH-NSi-16-14Fig. 1.— The rest-mass density profiles on the meridional plane for the NS-NS (SLy, Mtot = 2.7M! , Q = 1.0) and BH-NS (H4, Q =3, χ = 0.75) models at 8.8 ms after the onset of the merger.

The density and Kyutoku+13
velocity structures of the ejecta
morphology of the ejecta depend sensitively on the proBH-NS
of compact binary mergers are determined bygenitor models. Inthe following, we briefly summarize(M. Shibata’stalk)
-12
NS-NS
the numerical-relativity
simulation
using SACRA
theseproperties of the NS-NS and BH-NSejecta.
code (Yamamoto et al. 2008). We follow the dynamNS-NS ejecta. As shown in Fig. 1, the NS-NS ejecta
ical ejecta with the numerical-relativity simulation
have spheroidal shape rather than a torus or a disk irreuntil the head of the ejecta reaches ! 1000 km (seespective of Q and EOS as long as a hypermassive neuHotokezaka et al. 2013 and Kyutoku et al. 2013 fortron-10
star is formed after the merger. The reason is asdetails). After that, the density and velocity structures
follows. The
origin ofthe ejectafor NS-NSmergers
can-101
2
3of the ejecta are modeled assuming homologous expanbe divided into two parts: the contact interface of twosion (Rosswog et al. 2013a). For the simulations, weneutron stars at the collision and the tidal tails formedi - zAt the contact inemploy a piecewise polytropic EOS with which the cold
during an early stage of the merger.
EOSs of neutron-star matter are well fitted (Read et al.
terface, the kinetic energy of the approaching velocities
2009). For systematicof the dependenceof mass Wanajo,of the 4twoKiuchi,
stars is converted
into thermalenergy throughMT, studiesHotokezaka,
Kyutoku,
Sekiguchi,
Shibataejectionon the(SLy,coldMtot
EOSsof !neutron-starmatter,for the NS-NS= 2.7M, Q = 1.0) andBH-NS we(H4, Q =shock heating. The heated matter at the contact inter2013,
ApJ,
inpressface(arXiv:1310.2774)
consider five cold EOSs: APR4 (Akmal et al. 1998)
expands into low-density region. As a result, thebluer

4redder

2Berger, Fong, & Chornock

Application to GRB 130603B
Berger+13
Tanvir+13
A. Fruchter’s talk
F IG . 1.— Images of the afterglow and host galaxy of GRB 130603B. Left: Magellan/IMACS r-band image at 8.1 hr, with the bright afterglow marked. Middle:HST/ACS/F606W image at 9.4 days. Right: HST/WFC3/F160W image at 9.4 days. In all three panels the circle marks the position of GRB 130603B, with aradius of 10 times the rms of the astrometric match between the Magellan and HST images (1σ ≈ 34 mas). A source coincident with the GRB position is clearly 5visible in the WFC3/F160W image, with no corresponding counterpart in the ACS/F606W image.

Hotokezaka+13

BH-NStainty is 34 mas (1σ). The WFC3/F160W observations reveal
Golenetskii et al. 2013). The spectral lags are 0.6 ± 0.7
a point source coincident with the afterglow position, with noms (15 − 25 to 50 − 100 keV) and −2.5 ± 0.7 ms (25 − 50SLy(M=0.02)
in the ACS/F606W
observation (Figto 100 − 350 keV),
ej and there is no evidence for extended20 corresponding sourceMS1(Mej=0.07)
20ure 1).
emission(Norrisal. 2013). The combination of theseH4(M=0.004)
H4(Mej=0.05)
ejet
To measure the APR4(Mbrightness of thesource we use pointproperties indicates that GRB 130603B is a short-hard burst.
=0.01)
rej
spread-function (PSF) photometry with the standard PSFSwift/X-ray Telescope (XRT) observations commencedr We model
22
about 59 s after the burstHand led to the identification22 fitting routines in the IRAF daophot package.
the PSF in each image using 4 bright stars toH
a radius of 0.85"" ,of a fading source, with a UVOT-enhanced position ofand apply the WFC3/F160W PSF to the point source using a
RA=11h 28m 48.15s, Dec=+17◦04" 16.9"" (1.4"" radius, 90%0.15"" radius aperture and a background annulus immediatelycontainment; Evans et al. 2013).
24Ground-based observations starting at about 2.7 hr re-24 surrounding the position of the point source to account for
the raised background level from the host galaxy. Using thevealed a point source slightly offset from a galaxy
tabulated zeropoint, we obtain mF160W = 25.8 ± 0.2 AB mag.
visible in Sloan Digital Sky Survey (SDSS) images(Levan et al. 2013; de Ugarte Postigo et al. 2013; Foley et al.26 To determine the limit at the corresponding position in the26ACS/F606W observation, we use the PSF to add fake sources
2013; Sánchez-Ramírez et al. 2013; Cucchiara et al. 2013a).
of varying magnitudes at the afterglow position with the IRAF
The point source subsequently faded away indicating that it
addstar routine, leading to a 3σ limit of mF606W ! 27.5 AB
is the afterglow of GRB 130603B (Cucchiara et al. 2013b).
2828mag (see also Tanvir et al. 2013).
Spectroscopy of the host galaxy and afterglow revealed acommon redshift of z = 0.356 (Thöne et al. 2013; Foley et al.
2013; Sánchez-Ramírez et al. 2013; Cucchiara et al. 2013a;
3. A COINCIDENT RED COUNTERPART: AFTERGLOW ORXuetal.
2013).
We
obtained
twosetsofr-band
observations
KILONOVA?
3030of GRB 130603B with the Inamori Magellan Areal Cameraof the point10
source in the0.1
1
10 on the Magellan/Baade 0.1Given the precise alignment
1and Spectrograph(IMACS) mounted
WFC3/F160W image with the optical afterglow position we6.5-m telescope on June 3.996 and 4.992 UT, and detectedconsider it to be days
related2 afterto GRBGRB
130603B.
The simplest
Rest-frame
daysafterGRB
130603B
Rest-frame
130603B
the fading afterglow using digital image subtraction with theexplanation is that the source is the fading afterglow itISIS software package (Alard 2000). The afterglow posi-

Magnitude (AB)Magnitude (AB)NS-NSConstraints on Progenitors/EOSs
Efficient mass ejection
Mej ~ 0.02 Msun
NS-NS
BH-NS
=> soft EOS=> stiff EOSMsun)

APR4
SLy
ALF2
H4MS1

11.1 km

11.4 kmR1.35
12.4 km13.6 km0.1

14.4 km
0.01
0.13

0.14

0.15

0.16 0.17
Mtot/2R1.35

0.18

0.19

0.2

Compactness (M/2R)

Mej (10Mej (10soft

NS-NS models

stiff
10soft

BH-NS models

14.4 km13.6 km1
12.4 km

-2

1

stiff

Mej/10-2
Msun

10-2

Mej/10-2
MsunMsun)

4

0.1

0.01
0.13

APR4
ALF2
H4
MS1
ALF2(7,0.5)
H4(7,0.5)
MS1(7,0.5)

0.14

0.15

11.1 km

0.16
0.17
MNS/R1.35

0.18

0.19

Compactness (M/2R)

Fig. 2.— Ejecta masses as a function of the compactness of the neutron star, which is defined by GMtot /2R1.35 c2 and GMNS /R1.35 c2
Hotokezaka,
Kyutoku,
Sekiguchi,
Wanajofor NS-NS and BH-NSmodels, respectively.
Left panel:MT,
NS-NSKiuchi,models. Eachpoint shows theShibata,ejecta mass forthe equal mass cases. Error
bars denote the dispersion of the ejecta masses due tothe various
mass778,
ratios.L16
Right panel: BH-NS models. The filled and open symbols
2013,
ApJ,
correspond to the models with (Q, χ) = (3–7, 0.75) and (7, 0.5), respectively. The blue shaded region in each panel shows the allowed ejecta

Summary• EM emission from NS-NS/BH-NS mergers• Accurate localization of GW sources• “Full” radiative transfer simulations• Higher opacity than Fe by a factor of 100• SED peak at near-IR• Bluer emission from BH-NS mergers• Observing strategy• 22-25 mag (i band) => 4-8m class telescopes• Extremely broad-line spectra• GRB 130603B• Constraints on progenitors and EOSs