Multi-messenger Astrophysics ofcompact binary and core collapse
Kenta Hotokezaka(RESCEU, U of Tokyo)
Collaborators:
E. Nakar (Tel Aviv), T. Piran (Hebrew), S. Nissanke (GRAPPA Amsterdam), K. Masuda(Osaka), A. T. Deller (Swinburne)
G. Hallinan, K. Mooley, M. M. Kasliwal, P. Beniamini (Caltech),
M. Shibata (AEI/YITP), M. Tanaka (Tohoku U.), D. Kato (Kyusyu U.),
H. Nagakura (Princeton)

CT)_NA_TU_R_E_V_O_L_.3_3_1_4_FE_B_R_U_A_RY_1_98_8_ _ _ _ _ _ _ _LETIERSTO NATURE

1988AJ.

Table VIL Averaged £/2? FT?/photometry for the early days of 1987A.

The first multi-messenger Astrophysics:
--SN
1987A
--LETTERS(i) 1987 AUG 1 - (2)OCT 31

4847 keV

DATE(UT)

TIHE

JD
2446800+--V (B-V) (U-B) (V-R) (R-I)

(3)

(4)

(5)

(6)

(7)

I

Vl

0.04

(8)
Note

Vl

-+C:
0.02
LETIERSTO
NATURE
NZ
50.954 4.81 0.09 -0.84 0.22 -0.03::J

3 3_3_1_4_FE_B_R_U_A_RY_1_98_8_ _ _ _ _ _ _ __NA_TU_R_E_V_O_L_.

L REVIEW

6 ApRIL 1987

Feb 24.454
1.138
0
1.508 51.324 4.62 0.12 -0.75 0.24 -0.02
SAA01
Feb 24.824
0
Feb
1.850
0.13
0.27
-0.02
CTIOl'ESO'LCQ
25.16651.666
4.61
-0.68
0
---2
1987 AUG 1 - OCT 31
2.704 52.520 4.51 0.22 -0.54 0.32 0.01 CTI01fSAAQl,ES0
Feb 26.020I
847 keVVl>Q) 4
3.654 53.470 4.45 0.38 -0.21 0.35 I0.05
Feb 26.970
CTI02|SAAQ1,ESQ
0.04
C:
Vl
Q)
CTI02,SAADI,ESO
Feb 28.020
4.704 54.520 4.43 0.53 0.32 0.39 0.04
-0.02
~1-+SN 1_NA_TU_R_E_V_O_L_.3_3_1_4_FE_B_R_U_A_RY_1_98_8_ _ _ _ _ _ _ _ LETIERSTO NATURE
Vl
C:
5.644 55.460 4.46 0.68 0.71 0.42 -+Feb 28.960
CTIQ2,SAA01,ESQ
0.08
I
Vl

eventse lastors on

I

I

I

I

&

I

&

&

I

1

LU

0
0

ctron

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ngle
4

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I

I

I

1

I

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r

r

I

8~18
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Notes:
1987 AUG 1 - OCT 31
NZ, Horeno and Walker (1987)
847 keV
SAADI, Henzies et al. (1987)
LCD, Shelton and Ciaría (1987)
1985 JUL
26 -(1987)
OCT 25
ESO, Cristiani
et al.
CTIOl, HattheMS (1987)
CTI02, this work01.

COS 8(e, LMC)

FIG. 3. Scatter plot of the detected electron energy in
Hirata et al. 1987,
seecosine
also Biontaal. 1987,
megaelectronvoltsand theof the et
between theangle
measured electron direction and the direction of the Large
Magellanic Cloud. The number to the left of each entry is thetime-sequential
event number from Table I. The two projections of the scatter plot are also displayed.

0.02

::J
0C: -0.04-----_J
0

I

Vl
Vl
C:Vl

0
0.04

-+Q) -0.02
0.02
-+-C:
Fig. C:::J

0

60

SN 1
3 The fitted 847-keV
0
3
for test
positions differin
I
a:::
Q)
0
Vl
Also shoright
C: ascension.
-0.04
0
0
---2
_J
that
the
line
emission
orig
Vl
-+-I
2
of the supern
Vl
C:>
C:
::JQ) flux. After that, the IUE flux could be ignored as it
average monochromatic
flux at 3100 ÂvioletQ) -0.02
0~ 1
60
0
4
-+-the data up to Febru
points,which
arechangedthe total optical plus IR
flux
by
less
than
1%.
DurSN 101985 JUL 26 - OCT 25
C:
fit by a straight line.
We have fit a line
ing those first eight days we must add the IUE flux, as well as
I
Q)
Q)
part of the Fig.
flux 3curveinterpolate
theTheandfitted847-keV-+-Vlflux between 3200 and 3600 A, which falls between the
C: -0.04
00 3andmonochromatic
flux
at
3100
Â differingto the ti
IUE
optical data, to the optical and IR flux. Since thefortestpositions_Ja:::Vlvations. Notethatascension.
the first opticaldatasupernova
was changing so fast, we must interpolate the IUE
-+- 0
Also
show
right
the bolometriccurvesare at origi
FebrdataC:
thatlightthelineemission::J to the time of the optical observations given in Tables
I
0.04
0.43 days before thefirst
data60po
VI 0and20.5
Vl
1.5 IUE flux and the2the0 superno
ofIUE
4VII. Figure 5 shows the dereddened
0
1985
JUL
26
OCT
25
extrapolate the IUE data to this date. T
E(MeV)
Vl
shown in Fig. 5 gives
us confidence th-+Q)
Fig. 3C The
0.02fitted 847-keV
Fig. 2 -+-The background-subtracted spectrum for SN1987A,will not be too
faroff.
After 25.0 Febr
0 3::J positions
for thetest
differin
accumulated
from 1 August to 31 October 1987. The line profilesflux dominateda:::
0 IUE flux, and by Fe
Also show
for the two 56 Co lines plus a power-law continuum are plotted asflux was onlyright
2% 0ascension.
that of the optical/IR
that the lineluminosities
emissionorigi
0a solid line.
Also shown is an equivalent spectrum accumulated The final bolometric
are
02
the
supernoofin 1985. In both spectra residual features can be seen near 1.17and are plotted in Fig. 6. At the same ti
Vl0.04
and 1.33 Me
V due to imperfect subtraction1.5
oflines from the2internalmonochromatic Ifluxes
based on the bro
C0.5
VlQ) -0.02
residual
atmosphericcontinuum
calibration source. Note also the
to
fit
the
data
to
a
functionwiththe s-+SN 1987A
E(MeV)VlC:
in both spectra.
Since the data are-+not fit too well by a bla
CQ) 0.02
Fig. 2 The background-subtracted spectrum for SN1987A,included all the points
in the fit. For th
C -0.04
accumulated from 1 August to 31 October 1987. The line profileswe have used the 0::JIUE
and optical photo
0 56 the line comes from the supernova. The quoted
assumingthat_J
0 have used the V tofor the two
Co lines plus a power-law continuum are plotted asfits. After that, we
error
in
the
flux
has
been
increasedoverthestatistical
error
to
I
0
0.04
a solid line. Also shown is an equivalent spectrum accumulatedfits.
Vl2
includeestimateof possible
to theThe overall bolometricin 1985.anIn0.5
both spectra
residual systematicfeatures1.5
cancontributions
be seen near
1.17
light
-90. curve-60.
Vl
E(MeV) oflines from the internal1987A underwent
inand 1.33 Me
V due
to imperfectmeasuredline
intensity.
Thissubtraction
estimateis based on backgroundCVl a rapid decreaseDe
-+Q)
-0.02
C
0.02
first
seven
days
after
core
collapse.
The
also
the
residual
atmosphericcontinuum
calibration
source.
Note
Fig.
2
The
background-subtracted
spectrumfor
SN1987A,
subtracted spectra accumulated during times of similar back-+SN 1987A
by a factor of
5 toC:40::Ja broad
maximumoin
both
spectra.
accumulated
from 1 August
to 31
October
1987. The
lineTheprofilesFig.
The
fitted847-keVgroundand occultation
conditionsinprevious
years.
spec560
thendimmedrapidly,
and
on
days
12
for
the
two
Co
lines
plus
a
power-lawcontinuumareplotted
as
Q)
for
test
positionsdifferintrum for 1985 in Fig. 2 is one such accumulation. None of these-0.040have enteredanC exponential
declina solidline.any
Also shown is net
an equivalent
accumulated
declination.
-1Also shown aspectrashow
flux theat 847spectrum
keV. Since
no quoted
systemassuming
thesignificant
line
comes
from
_J day
mag
.
The
color
in 1985.that
In bothspectraresidualfeaturessupernova.
can be seenThenear 1.170.010 + 0.002
the lineemission originatev
Vlaticand847-keVline
is toobserved
insubtraction
theseoverpre-supernova
spectra,weerror
in1.33
theMe
flux
has
been
increased
the statisticalerrortheto supernova
declinedin temperatur
V due
imperfectoflinesfrom theinternal
Csupernova
1.source.
(a,b)
OpticalUBV{RI)kc
light
curvesforcontributions
1987A
in theoffittedfluxesas anestimateof 7 days to about
have
used
observed
rangeQ) -0.02
first
5500
K and -60.
has r
includeanFig.theestimateofpossiblesystematicto thealso
theresidual
atmosphericcontinuum
calibration
Note-+-90.
SN 1987A
LMC.
The
photometryisplottedasafunction
of
time
in
days
since
the
possiblesystematic
contribution
to
the
line
intensity
in
1987.
near
that
temperature
for
150
days.
T
C:
in bothspectra.
measured thelineintensity.
Thisestimatebased onbackgroundneutrino
outburst
detected
at IMBand isKamiokande
II detecDe
Matz
et
al.
1987
and
Mario-Hurmy
et
al.
1988
slightcoolingafter
themaximum
at
da
This
is
added
to
the
statistical
error
from
the
fit
to
give
an
backsubtractedtorsspectra
Q)(Koshiba accumulated
et al. 1987; Svobodaduring
et al. 1987).times
The firstof20similardays
nosity over Fig.
the firstdays
6.4 X10C 177-0.04
estimatederror
the inferred
photonagain
aretotalplottedin (a) aton
anconditions
expandedscale.in
Theprevious
data for the flux.
first weekNote
4
The
fittedis847-keV
groundandoccultation
years.
The
specassumingarethat
the linecomesgiven
from
theVII,
supernova.
The quoted_J
Once wehaveestimates
for
the
lum
the
averagedphotometry
in
Table
and
the
remainthe
847-keVline
as
data
ta
that
this
error
does
not
reflectthestatisticalsignificanceoftheforwetest
differingtrum
in Fig.
2been
is onesuch accumulation.
None ofthese
errorforin 1985
the
flux
has
increasedover
the statisticalerror
totemperature,
canpositions
estimate
the
phot
27
ing
data
are
the
nightly
averagesof
CTIO
photometry.
We
have
at
least
two
Mg
half-liv
Also
shown
declination.
line
detectionitself,which
isgreaterthan
5a.
spectraany significant
flux
atreported
847 keV.
Since no systemand radii are given ai
also estimate
includedthe ofearliest
Vnet
magnitudes
byMcNaught
includeshowanpossiblesystematic
contributions
to final
the temperatures
theline
originate
-90.
-60.
no
correlation
of
line
inte
The
spectrumandthesignificanceoftheincrease
in
line
atic
847-keVline
isZoltowski
observedthese pre-supernova
spectra, ples
we of the blackbodyemission( 1987)
andetThisal. (in1987).
fits areshown
inmeasuredline
intensity.
estimateis based on backgroundsupernova
De
betweensourceandbackintensity
indicatethat theobservedfeatureisasunlikely
to beofa
offittedfluxesanestimate
have
used
the
observed
range
subtracted spectra accumulated during times of similar backstatisticalfluctuation. contribution
The possibility theremains
that it isrealnova. All data used wethepossibleline intensity
in a1987.
Fig. 4 The fitted 847-keV
groundandsystematic
occultation conditions toin previous
years. The
specline produced
locally,either in error
the Earth's
atmosphereor anbysignificant
SAA passage,
This
is
added
to
the
statistical
from
the
fit
to
give
for test positions
differingtrum for 1985 in Fig. 2 is one such accumulation. None of these
activeswithsubstantially
activation totalof spacecraftmaterial.
There
is no flux.
observed
atmosestimated
error
on
the
inferred
photon
Note
again
Also
shown
a
declination.
spectra ©
show
anysignificant
net
flux
at
847
keV.
Since
no
systemAmerican Astronomical Society • Provided by the NASA Astrophysic

---

γ-ray

~ 30

0+ 30

data

I

7~

8+ 32

~1
Q)

1

grees)

4

0

LU

4

>Q)00

-+-

I

Q)

Vl

I ::J

I

C:

0

-+---2C:Z

Neutrino

09, 21I

UJ

Gauss-

y.

I

3

UBVRI

This multi-messenger observation confirmed that the birth of a neutron staris responsible for a supernova. (I’ll come back to this at the end of the talk)Gravitational-wave Astronomy

PRL 116, 061102 (2016)

PHYSICAL REVIEW LETTERS

properties of space-time in the strong-field, high-velocity
regime and confirm predictions of general relativity for the
nonlinear dynamics of highly disturbed black holes.

week ending
12 FEBRUARY 2016

the coincident signal GW150914 shown in Fig. 1. The initialdetection was made by low-latency searches for genericgravitational-wave transients [41] and was reported within
three minutes of data acquisition [43]. Subsequently,
matched-filter analyses that use relativistic models of compact binary waveforms [44] recovered GW150914 as themost significant event from each detector for the observations reported here. Occurring within the 10-ms intersite

Binary Black
Hole (BBH) merger
II. OBSERVATIONOn September 14, 2015 at 09:50:45 UTC, the LIGO
GW150914
Hanford, WA, and Livingston, LA, observatories detected

Binary Neutron Star (BNS) merger
GW170817

The Astrophysical Journal Letters, 848:L12 (59pp), 2017 October 20

Abbot

Figure 2. Timeline of the discovery of GW170817, GRB 170817A, SSS17a/AT 2017gfo, and the follow-up observations are shown by messenger and wave
gravitational-wave
event. Two types of information are shown for each band/messenger. First, the shaded dashes represent the times
the timepanels)
tc of theand
FIG. 1. The gravitational-wave event GW150914 observed by the LIGO Hanfordrelative(H1, lefttocolumn
Livingston (L1, rightinformation
reported in
GCN
Circular.
The names of the relevant instruments, facilities, or observing teams are collected at the beginning of the row. S
column panels) detectors. Times are shown relative to September 14, 2015 at 09:50:45
UTC. For was
visualization,
allatime
seriesare filteredrepresentative
observations
Table
in each band are shown as solid circles with their areas approximately scaled by brightness; the solid lines indicate whwith a 35–350 Hz bandpass filter to suppress large fluctuations outside the detectors’
most sensitivefrequency(seeband,
and1)band-rejectsourcewasdetectablebyTop
at least
telescope.
Magniﬁcation insets give a picture of the ﬁrst detections in the gravitational-wave, gamma-ray, optical, X-ra
filters to remove the strong instrumental spectral lines seen in the Fig. 3 spectra. Top
row,
left:
H1 strain.
row,oneright:
L1 strain.
þ0.5
They
respectively
GW150914 arrived first at L1 and 6.9−0.4
ms later at H1; for a visual comparison, radiothe H1bands.
data arealsoareshown,shifted illustrated
in time byby
thisthe combined spectrogram of the signals received by LIGO-Hanford and LIGO-Livingston (see Section 2.
Fermi-GBM
and INTEGRAL/SPI-ACS
amount and inverted (to account for the detectors’ relative orientations). Second row:
Gravitational-wave
strain projected lightcurves
onto each matched in time resolution and phase (see Section 2.2), 1 5×1 5 postage stamps extracted from the iniof SSS17a/AT
four
early spectra taken with the SALT (at tc+1.2 days; Buckley et al. 2017; McCully et al. 2017b), ESO-N
detector in the 35–350 Hz band. Solid lines show a numerical relativity waveform observations
for a system withparameters 2017gfo
consistentand
with
those
+1.4days;
Smarttetal.
2017),
the
SOAR
4
m
t
c based on [15]. Shaded areas show 90% credibletelescope (at tc+1.4 days; Nicholl et al. 2017d), and ESO-VLT-XShooter (at tc+2.4 days; Smartt et al. 20
recovered from GW150914 [37,38] confirmed to 99.9% by an independent calculation
described
in binary
Sectionblack
2.3, holeand theﬁrst X-ray
and radio detections of the same source by Chandra (see Section 3.3) and JVLA (see Section 3.4). In order toregions for two independent waveform reconstructions. One (dark gray) models the signal using
template
waveforms
representative
spectralenergydistributions,each
spectrum is normalized to its maximum and shifted arbitrarily along the linear y-axis (no absolute scale). Th[39]. The other (light gray) does not use an astrophysical model, but instead calculates the strain signal as a linear combination of

The first GW detection
Abbott et al. 2016

The first GW & photon detection
Abbott et al. 2017

Outline•

A comment on Binary Black Holes: an astrophysicalimplication from new LIGO/Virgo catalogue (GWTC-2)•

Binary Neutron Star: what we have learned from multimessenger observations of GW170817•

High energy (~100 MeV) neutrinos from core collapse
supernovae and prospects with HyperKamiokande.

I will discuss very high energy gamma-rays from mergers in the
CTA meeting next week.

Gravitational-Wave Transient Catalogue 2 (GWTC-2)
GWTC-2 contains 50 GW events detected by LIGO and Virgo since92015 to O3a (Abbott et al 2020).

ate events for the first timeeported candidates. Among
gravitational wave emission
ence of BBHs, BNSs, andCumulative Detections

our use of multiple search70
ate events, we expect ⇠ 3vents caused by instrumen60
his catalog. It is not posainty which specific candiO3a
50 O1 O2e; instead we provide staarm rate and probability of
40
these candidate events, 26
y in real-time processing via
3029]. Furthermore, four gravnts from O3a have already20
ue to their interesting prope second gravitational wave
10
coalescence; GW190412 [30]with definitively asymmeth also produced detectable
0
0
500
1000
1500
2000
nd the leading quadrupolareven more asymmetric sysE↵ective BNS VT [Mpc3 kyr]t merging with a ⇠2.6 M
her the lightest black hole or
to be in a double compact
FIG. 1.
The number of compact binary coalescence de32, 33] isBBH
a BBHmerger
with totalrate:tections
versus the e↵ective +44(VT) to which the+138 65-1 (2016)
mary massat -3 yr
99above
56 volume-time
23.9 +14.9 -8.6 Gpc-3 yr-1 (2020)-70MGpc-27 (2018)
gravitational wave network is sensitive to BNS coalescences.
The e↵ective VT is defined as the Euclidean sensitive volume~ 2.4 Myr-1 in the Milky Way[34] of the second-most sensitive detector in the network at agiven time, multiplied by the live time of that network conweresensitive
luckyvolume
for BBHs.
figuration.=>
TheWe
Euclideanof each detector
is calculated from the BNS inspiral range [34, 35] shown inSpin distribution indicates the field binary origin
Among 46 BBH mergers, 11 mergers have a positive χ within
a 90% credible interval. On the contrary, there is no event with aclear negative χ

ack hole spin magnitude and tilt distributions. Left: The distribution of dimensionlesseff
Default spin model (see Appendix D.1). Light traces show individual draws from thek curve shows the posterior population distribution for . Dashed lines mark the central
d distribution of tilt angle cos ✓1,2 of black hole component spins relative to their orbital
rientation, which corresponds to a uniform
in therecos ✓1,2is, isadisfavoredbutχeff
not large observable volume)(Note thatbias, largeeff. distribution
out a highly peaked distribution at cos ✓1,2 = 1. Rather, the data are consistent with aest preference for aligned spin (cos ✓1,2 > 0).

In fact, the underlying spin distribution peaks at positive χeff.
6
Default
Gaussian5(Abbott et al 2020)

p( eff )

4
3
2
1
0

0.6

0.4

0.2

0.0

0.2

0.4

0.6

eff

mean µe↵ and standard deviation e↵ of the BBH e↵ distribution, obtained using theD.2. We
marginalize
over
the parameters
the angular
distribution momenta
of eﬀective precessing
Positiveχeff
means
that governing
the spinand the orbital angular momentum
n that is peaked at positive values, its measured width implies that a non-zero fraction ofare aligned, which is a very strong expectation of the field binary scenario.
ying component spins misaligned by t > 90 relative to the orbital angular momentum.
ns for the eﬀective inspiral spin e↵ obtained with both the Gaussian and Default spinl 90% credible bounds on p( e↵ ) and the solid lines show the median posterior prediction

Outline•

A comment on Binary Black Holes: an astrophysicalimplication from new LIGO/Virgo catalogue (GWTC-2)•

Binary Neutron Star: what we have learned from multimessenger observations of GW170817•

High energy (~100 MeV) neutrinos from core collapse
supernovae and prospects with HyperKamiokande.

Neutron Star Merger
Propositions

Since GW170817•

One of the strongest GW emitters (Hulse & Taylor 1975)•

Standard cosmic siren (Schutz 1986)•

A site of heavy nucleosynthesis (Lattimer & Schramm 1974)•

Short GRB progenitors (Eichler et al 1988)•

A laboratory of high-dense material (Flanagan & Hinderer 2008)•

Fast Radio Burst progenitors (Totani 2013)?

EOS=APR, Mtot = 2.7Msun

EOS=APR, Mtot = 2.9Msun

KH + 13

Variety in merger remnants

Figure 1

Shibata & KH 2019

A summary for the merger and post-merger evolution of binary neutron stars. Mthr andThe first binary neutron star merger:
GW170817
PHYSICAL REVIEW LETTERSEVIEW LETTERS

week ending
20 OCTOBER 2017

PRL 119, 161101 (2017)High Signal-to-noise ratio ~ 30

TABLE I. Source properties for GW170817: we give ranges encompassing the 90% credible int
waveform model to bound systematic uncertainty. The mass values are quoted in the frame of thethe source redshift.

1.36–1.60 M ⊙1.17–1.36 M ⊙1.188þ0.004−0.002 M ⊙0.7–1.0
2.74þ0.04−0.01 M ⊙> 0.025M ⊙ c2
40þ8
−14 Mpc
≤ 55°
≤ 28°
≤ 800
≤ 800

BH

lim
it

Primary mass m1
Secondary mass m2
Chirp mass M
Mass ratio1m2 =m1
Total mass mtot
Radiated energy Erad
Luminosity distance DL
0 Θ
Viewing3.0
angle
Using NGC 4993 location it
lim deformability Λ~
Combined dimensionlessl tidalhDimensionlessΛð1.4M ⊙ Þda
2.5 tidal deformability

Low-spin priors ðjχj ≤ 0.05Þ

chu
B

FIG. 1. Time-frequency representations [65] of data containing
the gravitational-wave event GW170817, observed by the LIGOHanford (top), LIGO-Livingston (middle), and Virgo (bottom)

Abbott et al 2017

H4

SLy
MPA1

APR4

WFF1

m (M )

to the one observed at the LIGO-Livingston detector during
From the gravitational2.0
GW170817. After applying the glitch subtraction techin the frequency range cnique, we found that the bias in recovered parameters
mass in the detector fram
relative to their known values was well within their[51]. The mass parameter
1.5 This can be understood by noting that a
uncertainties.
to the rest-frame masses
small time cut out of the coherent integration of the phasemdet ¼ mð1 þ zÞ [142].
evolution has little impact on the recovered parameters. To[90], and correcting forcorroborate
theseresults,
the
test
was
also
repeatedwitha
1.0
Barycenter with respe
window function applied, as shown in Fig. 2 [73].
Background [143], the g
The source was localized to a region of the sky 28 deg2
urement alone implies
in area, and 380 Mpc3 in volume, near the southern end of
0.5
0.008þ0.002−0.003 , which is cothe constellation 8Hydra, by 10
using a combinationof14the 0
12
1[50,141,144,145].
Withotiming, phase, and amplitude ofRthe(km)source as observed inin the source’s chirp
the three detectors [138,139]. The third detector, Virgo, wasuncertainty in its lumino
essential in localizing the source to a single region of the
waveform model or the csky,
in Fig. 3.posterior
The smalla
FIG.as3.shown
Marginalized
forskythearea
masstriggeredm and areal
radius R of each binar
thesource-frame chirp m
successful
follow-up campaign
thatweidentified
electroand a parametrized
EOS where
imposeana lower
limit on the maximum mass
While the chirp mass
magnetic counterpart [50].

Merger rate
Gravitational wave
RGW =920+2220-790 Gpc-3 yr-1 (2018)=> ~ 32+49-24 Myr-1 the Milky-Way Galaxy

320+490-240 (2020)
We were also lucky for neutron star mergers.

Short GRBs
RSGRB = 6+2-2 Gpc-3 yr-1 (before a beaming correction, Wanderman & Piran 15)=> ~390+130-130 (fb-1/65) Gpc-3 yr-1 (corresponding to a half-opening angle of 10o)
RGW ~ RSGRB (corrected) suggests that all short GRBs can arise from mergers.

Galactic binary neutron stars (BNS)RBNS = 42+30-10 Myr-1 (Pol, MclLaughlin, & Lorimer 2019),
which is dominated by 5 systems, J1906+0746, B1913+16 (Hulse-Taylor),
J0737-3039A/B (the double), J1757-1854, J1946+2052 (the tightest)

EM Counterpart Discovery Timeline

Follow-up observations of GW170817
Gravitational waves(2017 Aug 17.5)
T=0Near Infrared
T = 11h 36m

X-rays
T = 9 days

Off-axis jet?

Delayed Arrival of X-rays and Radio

Gamma-rays
T = 1.7 seconds

OpticalT = 10h 52m

Ultraviolet
T = 15 hours

Radio
T = 16 days

Extremely faint!
(~15 uJy @ 3 GHz)
Abbott et al., ApJL. 848, L12 (2017)

the jet eventually burrowed through the ejecta to2018).
See
Figure1
foradepictionofthedi↵erent
GW170817: GRB, Kilonova & Afterglow
31

Goldstein +17

GRB 170817 (X-γ)

Kilonova (uv-IR)e 7. The 256 ms binned lightcurve of GRB 170817A in the 10–300 keV band for NaI 1, 2, and 5. The shaded regions are

Afterglow (radio-X)

Figure1:Leftpanel:
The
optical/NIR“kilonova
Dissipation in the outflow:
Kinetic energy deposited
Radioactive decay:
signalis poweredtherainto the
ISM: ~ 1038 by- 1040erg/s
42 erg/s
L ~ 1046 - 1047 erg/sThe short-lived~ 1038 - 10kilonova↵erent time intervals selected for spectral analysis. The inclusion of the lower energies shows the soft tail out to T0+2 s.

Kilonova in GW170817
Arcavi+17, Coulter+17, Lipunov+17, Soares-Santos+17, Tanvir+17,
Valenti+17, Kasliwal+17,Drout+17, Evans+17, Utsumi+17−19
7ype Ia
supernova
7ype Ib/c
supernova

Absolute magnitude

−18
−17

61 2002bj−16−15

61 2010X−14−13

61 1999em + 1

Arcavi + 17

Kilonova

−12
−20

−10
0
10
20
30
7ime Irom peaN (rest-Irame days)

Basics of Kilonovae
Li & Paczynski 1998, Kulkarni 2005, Metzger + 2010

MergerMass ejection

R-processnucelosynthesis

Radioactive decay
Photons
diffuse out

Optically
thin

Time
t=0

0<t<100ms

~<1s

1day

> 10day

10km

10-100km< 0.1Rsun

10AU>100AUDynamical mass ejection
1.5Msun and 1.2Msun

1.6Msun and 1.3Msun

KH + 2013

HMNS formation:• Tidal (cold) + shock (hot)• Ejection lasts: ~ 5 ms
• Mass < 0.01Msun, v ~ 0.2c

Prompt BH formation:• Tidal (cold)• Ejection lasts: ~ 1 ms
• Mass < 0.01Msun, v ~ 0.2c

Bauswein + 13, Piran + 13, Rosswog 2013, Kyutoku+15, Sekiguchi + 15, 16, Radice+16

DynamicalejectionvsDisk
outgflow
16

Disk outflow
Mej;sec [M ]10 1

Radice, Perego, KH + 18 (see also Wu+16 and Shibata+17)

10 2
10 3
BHB⇤
DD2
LS220
SFHo

10 4
10 5 5
10

10 4

10 3
10 2
Mej [M ]Dynamical ejecta

10 1

Figure 16. Dynamical ejecta Mej;dyn versus secular ejecta masses
Mej;sec . With the exception of the prompt BH formation cases that

R-process nucleosynthesis in merger© Jonus Lippuner

2
r-processproduction
rate,
sGR
Kilonova Emission depends on the composition
Light

Cumulative mass abundance in the sun

3rd peakLanthanide

10-6

2nd peak-5

1st peakSolar r-process abundance

10

Heavy

10-7
10-8
10-9
-10

10

10-11

80

100

120

KH, Beniamini, Piran 18, Goriely 99

140

160

180

Atomic mass number

200

220

1

1

1

1

10

Heating rate of r-process
Way & Wigner 1948

KH & Nakar 2020

K. %AY AN l3 E- P. WI GNE R

1016

IQ-

«I

10

Heating rate [erg/g/s]

10

IO

«

lO

C-.

+

IF-

14

t-1.3

1012

10

10

108

g'

I

I

.OOI

10

.

I

.I

l

NI

SIONIOI

Q

lO

Q

i~

I6
I deyap

Is

KPdoye ICPOd~aHeating rate of nuclear waste

10erated.

I (I)FSNl g KNO

6

Fj:G. 6. Theoretical and experimental results for-2the radiotivity t seconds after fission.
of P-particles
P(t) = number
emitted per second, C7(t) samefor y-rays' B{t}=energy of Pparticles in Mev per second,F(t) =energy of y-rays in sameunits. 38(t)+I'(t) is total energy(including neutrino energy) lib-

10

10

0

2

10
10
Time [s]

4

10

6

10

Heating rate of r-process8Heating rate of r-process
Way & Wigner 1948
IQ-

«I

IO

«

lO

C-.

+

Heating rate [erg/g/s]

10

IF-

g'

KH & Nakar 2020

K. %AY AN l3 E- P. WI GNE R

1016
1014
1012
1010
108

I

I-2 .

.OOI
106 .

I

l

I

10

0

SIONIOI

10

Heating rate of nuclear waste
10

NI

2

Q

lO

10
10
Time [s]

4

Q

I6

Is

p ICPOd~ai~ 106 KPdoye
108
I deyaFj:G. 6. Theoretical and experimental results for the radiotivity t seconds after fission.
of P-particles
P(t) = number
emitted per second, C7(t) samefor y-rays' B{t}=energy of Pparticles in Mev per second,F(t) =energy of y-rays in sameunits. 38(t)+I'(t) is total energy(including neutrino energy) lib-

Heating rate of r-processerated.

I (I)This is somewhat a unique properties of the heating rates ofmany beta-decay chains.
FSNl g KNO

Observation vs theory of Kilonova
Open code: https://github.com/hotokezaka/HeatingRate
KH & Nakar 2020

Heating Rate

Lbol

~0.05 Msun of r-process elements are required to power thekilonova GW170817.

open s shell(l=1)

Radioactive heat => Photon Luminosity
Barnes & Kasen 13, Kasen + 13, Tanka & KH 13, Tanaka 17, Wollaeger + 18, Tanaka+19

Photons are blocked by atomic transitions

Expanding Ejecta

τ<1open s shell(l=1) tdiff<t

5

5

4

4

open p-shell3(l=2)

E (eV)

open d-shell
(l=3)

5

2

2

1

1

0

4

γ

E (eV)

3

0

Fe II

Nd II

Nd II

0

open f shellR-process
Opacity (l=4)

 ⇠ 0.1 cm2 /g

2

1

E (eV)

3

tdiff>topen d(l=

10 cm2 /g

Kilonova GW170817: R-process heating matches the data
open s shell(l=1)

Open code: https://github.com/hotokezaka/HeatingRate

Expanding Ejecta: 0.01Msun , 0.1c
Radioactive(l=1)power
Iron-like

Tanaka& KH 13 (see also Barnes & Kasen 13)
open s shell

τ<1tdiff<t4

4open p-shell3(l=2)
E (eV)

E (eV)

Mixed
5

open d-shell
(l=3)

2

2

1

1

0

2

1

Fe II

Nd II

0

open f shell(l=4)
2
 ⇠ 0.1 cm /g
10 cm2 /g

3
E (eV)

tdiff>t5

3

4

Only Heavy element

5

Nd II

0

open f shell(l=4)

R-process mass
~ 0.05Msun

open d(l=R-process mass budget from GWTC-1
KH, Piran, Paul 15, KH, Beniamini, Piran 18

°1
°1
°1
Galactic
rate
[Myr
rate
[Myr
]]
Galactic

333
10
10
10

2
10
10
1022

1
10
10
1011

RRp
prroo
ccees
sssm
maas
ssin
inM
MWW

Prediction
After O3
in 2016
half

GW 170817

GW BNS

Galactic
BNS
Galactic
GalacticBNS
BNS
±± ±
sGRB
==10
10
sGRB
sGRBµµjµj j=
10

244
244Pu

Pu

Dwarf
Dwarf

AAmin
=69
69
min
min =
0
10
10
1000 °3
°3
10
10
10°3

°2
°1
°2
10
10
10
10°1
Mass
Mass perperevent
event[M[MØØ] ]

000
10
10
10

Ref: Goriely 1999, Lodders et al 2009, Wanderman & Piran 2015, Fong+2015, KH, Piran, Paul 2015,
Beniamini, KH, Piran 2016, Pol, McLaughlin, Lorimer 2019, KH & Nakar 2020, LVC 2020

R-process mass budget from GWTC-2
KH, Piran, Paul 15, KH, Beniamini, Piran 18°1
Galactic rate
rate [Myr[Myr°1
]]

33
10
10

Rp
ro c
ess

ma
ss
in

10
1022

Prediction
After O3
in 2016half
MW

GW BNS

Galactic BNS
BNSGalacticsGRB µµjj =
= 10
10±±
sGRB

10
1011

244
Pu

Pu

Dwarf

Amin
min = 69
101000 °3
10
10°3

10°2
10°1
Mass per event [MØ ]

1000
10

Ref: Goriely 1999, Lodders et al 2009, Wanderman & Piran 2015, Fong+2015, KH, Piran, Paul 2015,
Beniamini, KH, Piran 2016, Pol, McLaughlin, Lorimer 2019, KH & Nakar 2020, LVC 2020

Sr lines in the kilonova spectrum
60
Sr ii

F [10 17 erg s 1 cm 2 Å ] + offset

50

40

Sr ii

Blackbody + Sr ii
Blackbodyblueshiftdecades later30 . However, no spectroscopic confirmation of the composition of neutron stars has ever been made. The identification here ofan element that could only have been synthesised so quickly under an
extreme neutron flux, provides the first direct spectroscopic evidence
that neutron stars comprise neutron-rich matter.

X-shooter spectra are explained byblackbody + Sr II lines.

Tanaka & KH 13

8

Tanaka & Hotokezaka

Sr II

Watson et al 19

inosity (erg s-1)

1

Residual F [10 17 erg s 1 cm 2 Å ]Log flux (Fλ) + constant1

1. Burbidge, E. M., Burbidge, G. R., Fowler, W. A. & Hoyle, F. Synthesis of
6
the Elements
in Stars. Rev. Mod. Phys. 29, 547–650 (1957).
2. Siegel, D. M., Barnes, J. & Metzger, B. D. Collapsars1.5
as daysa major source
1.5 daysof r-process
5 elements. Nature 569, 241–244 (2019). 5.0 days
30 offsets
3. Lattimer, J. M., Mackie, F., Ravenhall, D. G. & Schramm,
D. N. The
10.0 days
decompression of cold neutron star matter. Astrophys. J. 213, 225–233
4
2.5 days(1977).
4. Eichler, D., Livio, M., Piran, T. & Schramm, D. N. Nucleosynthesis, neu20
3 and gamma-rays from coalescing neutron stars. Nature 340,
trino bursts
3.5 days
126–128 (1989).
5. Freiburghaus,C., Rosswog, S. & Thielemann, F.-K. R-Process in Neutron
2
Star Mergers. Astrophys. J. Lett. 525, L121–L124 (1999).
10
6. Ji, A. P., Frebel, A., Simon, J. D. & Chiti, A. Complete Element Abun4.5 days
1 Nine Stars in the r-process Galaxy Reticulum II. Astrophys. J.
dances of
10000
15000
20000
830, 93 (2016). 5000
7. Metzger, B. D. et al. Electromagnetic
counterpartsof compact object
Wavelength(A)mergers powered by the radioactive decay of r-process nuclei. Mon. Not.
3
1.5 daysFig. 6.—UVOIR spectra
model NSM-all
at 2650–2662
t = 1.5, 5.0, and10.0 days after the merger. The spectra are almos
R.ofAstron.
Soc. 406,(2010).
-3
wavelengths, there are possible absorption troughs, which result from Y i, Y ii, and Lu i in our simulations. Howev
8. Barnes, J. & Kasen, D. Effect of a High Opacity on the Light Curves ofbe result2.5from
daysthe incompleteness of the line list in the NIR wavelengths.
3Radioactively Powered Transients from Compact Object Mergers. Astro-3
phys. J. 775, 18 (2013).
3
3.5 days
9.
Tanvir,N. R. et al. A ‘kilonova’ associatedthe short-duration-ray of GW•
of withEM radiation
to the analysis-31042
burst GRB 130603B. Nature 500, 547–549(2013).
able
to pin down the masses of two NS
3
1.2 Abbott,+ 1.5 B. P., Abbott,
4.5 days
NSM-all
10.
R.
&
et
al.
GW170817:
Observationof Gravitationalof the EOSs moreaccurately. Note th•1.3 + 1.4-3
APR4(soft)
Waves from a Binary Neutron Star Inspiral.
PhysicalReview
Letters
119, per masssimulations, the heating rate
H4 (stiff)
4000
5500
7500
10500
14500
20000
161101 (2017).
understand the connection between th
11.
Pian,
E.
et
al.
Spectroscopic
identification
nucleosynthesis
of ofther-process
NS merger
and expected emissio
Observed wavelength [Å]
41
10
in a double neutron-star merger. Nature 551,
67–70
(2017).
osynthesis calculations are necessary.
12. Smartt, S. J. et al. A kilonova as the electromagnetic counterpart to a
Figure 4 | Spectral series of AT2017gfo 1.5–4.5 days after the merger. Data
gravitational-wave source. Nature 551, 75–79 (2017).
6. IMPLICATIONS FOR OBSE
are shown in grey and have been smoothed slightly. A model (solid red lines)

Sr lines are expected to be very strong.
Heavy elements may be absent in theouter part of the ejecta.

2
r-processproduction
rate,
sGR
Kilonova Emission depends on the compositionLight

3rd peakCumulative mass abundance in the sun

10-7

Lanthanide

10-6

2nd peak-5

1st peakSolar r-process abundance

10

Heavy

Sr

10-8
10-9
-10

10

10-11

80

100

120

KH, Beniamini, Piran 18, Goriely 99

140

160

180

Atomic mass number

200

220

1

1

1

1

10

Proposed Scenarios for the GW170817 kilonvoa
Most of mass: lanthanide-richlanthanide-less
Waxman + 2017, Shibata..KH+17,
Kawaguchi+18

Kasen +17, Villar + 2017
Perego+2017bcsquHHzHd dynamicalv  0.2c-0.3c

al dynamical
0.2c-0.3c
disN windv ≲ 0.1c

disN windv ≲ 0.1c

tidal dynamicalv  0.2c-0.3c

V ILLAR ET AL .

n Star1Hutron Star + 1Hutron Starmnant

rHmnant SromSt collaSsH to blacN holH

1Hutron Star + BlacN +olH
blacN 10holH rHmnant
Figure 6

Dynamical ejecta
4Analytic Model°°

AB absolute magnitude

Apparent Magnitude

D=40 Mpc, 20 )e<28
15
Schematic
picture of the ejecta profile for the case that a long-lived massive
neutron star is formed-17
as a remnant. The large spheroidal-shell component denotes the neutron-rich dynamical gejecta.
r
The small spheroidal-shell and polar components denote the-16 early viscous and long-term iviscousz
components, respectively. The “Low Ye ” implies that it contains neutron-rich matter withJ
H
-15 the nucleosynthesis of lanthanideYe <⇠ 0.2, which contributes to enhancing the opacity throughKelements. The “Medium Ye ” and “High Ye ” imply that it does not contain such neutron-rich
20-14
U
matter because
Ye >⇠ 0.25 and Ye >⇠ 0.40, respectively. The filled circle and ellipsoids in the central
Bregion denotea massive neutron star and accretion torus surrounding it. We note that the “Lowg
-13
Ye ” component
has high average expansion velocity of vej ⇠ 0.2c while the “Medium” and “High”i
J − 2have slower velocity, 0.1–0.2c. Note that the gravitational-wave observation indicatescomponents-12
H−2
that we observethemergerremnantof
GW170817
alongthedirection of ✓ <⇠ 30 from theK−2-11rotationaxis.
25

hematic illustration of the components of matter ejected from neutronRed colours denote regions of heavy r-process elements, which radiate

ght. Blue colours denote regions of light r-process elements which radiate

10 0

10 1

cess ejecta in the plane of the binary. Material squeezed into the polar
The variance parameter is an additional scatter term, which
we fit, that encompasses additional uncertainty in the models
and/or data. For upper limits, we use a one-sided Gaussianpenalty term.
For each component of our model there are four free parameters: ejecta mass (Mej ), ejecta velocity (vej ), opacity (),and the temperature floor (Tc ). We use flat priors for the firstman et al. 2014), which accounts for both the likelihood score
and number of fitted parameters for each model.

2

4. RESULTS OF THE KILONOVA MODELS/s/Å]Remnant:
lived
NS
=> BH
g the stellar
collision can shortform a coneof light
r-processmaterial. RoughlyLong lived neutron star (~100 ms)
We fit three different models to the data: a sphericaltwo-component model, a spherical three-component model,

-

-

-

-

-

0 1 2 3 4 5 6 7 8 9 10 11 12 13 14 15 16
F IG . 11.— A comparison of the line expansion opacity for Fe and Ce, att [day]Time since explosion [days]
T = 5000 K, ρ = 10−13 g cm−3 and t = 1 d, adopted from Kasen et al.
(2013),Fig.
using2.—the Autostructure,
and VALD oflineSSS17a
lists (Notethat sincewith the
Optical and Kurucz
NIR lightcurves
compared
F IG . 10.— A comparison of the observed multi-band light curves and thethe expansion
opacityis proportional
to 1/(ρct),
opacityfor our
86  ✓ 
90 (right
panel). The
optical theandimplied
NIR data
pointsare taken
analytic model multi-band light curves, given by Eqs. (12), (18) and (19), forinferredmodel
parameters
at
t
=
7
d
would
be
≈
7
times
largerthanof 40 Mpc. All the magnitudes are given in AB magnitudes.shown
Note that t
the11.
sameS.
model
parameters M.
usedLiebendoerfer,
in Fig. 9 (and including
extinction).
The
Rosswog,
F.-K.
Thielemann,
M.
B.incompleteness
Davies, W.ofBenz,
Astron.
toplot).
thethe line andlist forT.
thePiran,
opacity estimation.
in
thismodel’s spectral energy distributions are obtained assuming a black-body
spectrum.Astrophys.
This assumption341
breaks499
at t (1999).
> 7 d (the dotted line marks the
opacity of the Lanthanides is much higher at IR wavelengths,
time at which the ejecta becomes marginally optically thin, see Fig. 8). Thus,° (2000).
2−1
12.weS.
Rosswog,
M.B.
F.-K.
Thielemann,
Piran,
Astrophys.
360
171
κν T.
≈ 10
cm2 gAstron.
, see
11. This
impliesthat
Iron° groupt = Fig.
1.4[day],
D = 40Mpc, 20)e<28
while
expectthe model toprovideDavies,
a good estimateof thebolometric lu- andminosity at t > 7 d, we do not expect it to provide an accurate description
ï16elements
will
13. L.-X. Li and B. Paczhński, Astrophys. J. 507, L59(1998).
4×10not be able to produce the observed IR lumiof the spectra. The fact that the observed spectrum is not far from thermal at= 5150[K] Lanthanides’
nosity, while theLanthanides may. TThe
large
2
t∼
7 d implies
the opacity atG.
thisMartinez-Pinedo,
time is ∼ 0.1cm /g at 1 −S.
2µm
(see
14.
B. D.that
Metzger,
Darbha,E. Quataert,
3×10ï16A. Arcones, D. Kasen, T. Thomas,

ght. During the merger, tidal forces peel off tails of matter, forming a torus
Figure 1. Complete UVOIR light curves, along with the models with the highest likelihood scores. Solid lines represent the realizations ofhighest likelihood for each model, while shaded regions represent the 1 uncertainty ranges. For some bands there are multiple lines that
capture subtle differences between filters.

-

AB absolute magnitude

6

tidal dynamicalv  0.2c-0.3c

Kilonova Nebula
Early times

Late times (nebula)τ<1tdiff<tτ<1tdiff>tNebular phase:• Most of the ejecta can be seen. (inner parts have slower velocities)• Photon luminosity ~ heating rate
• Photons are emitted directly by radiative de-excitations.

Future Prospects of Kilonova

NeutronStar
MergerNeb
When the ejecta becomes optically thin, we
will see atomic
emissionlines.
10L⌫ [erg s 1 ]

10

39

[µm]

1

10

0

T = 40 day

Spitzer 4.5 μm (GW170817)

t = 40 day

t = 80

1038
JW

ST
M

IR

1037

1036 1310100 Mpc

c

IL

RS

T
WS

J

1014

[Hz]

N

pe
S
IR

1015 1013
KH et al in prep.

James Webb Space Telescope will be extremely powerful to get kilonova spectra!!

A Gamma-Ray Burst after GW170817
The Astrophysical Journal Letters, 848:L12 (59pp), 2017 October 20

Gravitational wave

γ-raysProperties of γ-rays:
Similar to normal GRBs
1) Delay is ~1.7 sec and duration is ~2 sec.
2) Isotropic energy is ~ 1047 erg and spectral peak is ~200 keV.
Much less than normal GRBs

3) Off-axis short GRB => spectral peak < 10 keV(e.g. Matsumoto+18)

On-axis mildly relativistic outflow
Explosion (merger) at t=0

Kasliwal…KH+17, Gottlieb, Nakar, Piran, KH 17
Also Beloborodov + 19

βsh, Γsh=(1-βsh2)-1/2

Shock

Rsh
1. Duration: Tobs ~Rsh/2Γsh2c ~ 1 sec. (observed)
2. γ-ray energy: E ~ Γsh Μc2 ~ 1047 erg (observed)
3. Optical depth: τ = κM/4πRsh2 ~1 (required)On-axis mildly relativistic outflow
Explosion (merger) at t=0

Kasliwal…KH+17, Gottlieb, Nakar, Piran, KH 17
Also Beloborodov + 19β, Γ=(1-β2)-1/2

Shock

R
1. Duration: Tobs ~Rsh/2Γsh2c ~ 1 sec. (observed)
2. γ-ray energy: E ~ Γsh Μc2 ~ 1047 erg (observed)
3. Optical depth: τ = κM/4πRsh2 ~1 (required)=> Rsh ~ 1011-1012 cm, Γsh ~ 3-5, Μ ~ 10-8-10-7 ΜsunTime delay: δΤ = (1-βej)Rsh/c ~ 1 sec => βej ~0.7 - 0.8
Merger simulations show a fast ejecta tail with ~0.8c and 10-7Msun (Kiuchi+17, KH+18)
But also see Ioka & Nakamura 2018 for off-axis jet considerations.

Late-time Afterglow across multi-wavelength
GW170817:
Afterglow
spectrum
Light curve (Makhathini+2020)
Spectrum

x10 7
x10 6

108

x10 5

37.6

27.5

Electron distribution:

37.4
•
28
Figure 1. Upper panel: The panchromatic (radio, optical and X-ray) afterglow light curve of GW170817, color coded according

dN
d e /

2.2

to the observing frequency, up to 940 days post-merger (all data points have 1 errorbars as presented in Table 1; upper limits
are not shown here) using the uniform dataset presented in this work. The light curve is scaled to 3 GHz using the best-fit
37.2
spectral index ( 0.569) derived
28.5 from the MCMC power-law fitting (see §3). Lower panel: The averaged (using moving average;
t/t = 1/17 where t is the width of the kernel and t is the time after merger) light curve (blue data points) shows a generaltrend consistent with power-law rise and decline. In grey are the same data points as shown in the upper panel.

e37

on the pn detector were 29
used to determine periods ofing to Chandra results from similar epochs. Backgroundhigh background data, which
were
excludedfrom
ourevents Merger
were extracted
from nearby, source-free circular
m
=
m
m
Time
after
(days)
1•
old in onlynew
analysis. Significant background flaring resulted
regions of 5500 in radius.
0 for obsIDs 0811210101
e were26 and 48 ks filtered exposures
For both GW170817 and NGC 4993, the spectraand 0811212701, respectively.
Sourceeventswere
thengroupedwithaminimum
of
1
count
per
bin
with
the
-1
extracted from a circular region1 with radius 500 centeredheasoft 2tool grppha. We used the X-ray spectral
3
10
10
10
on GW170817. While this is less than 50% of the EPICfitting package xspec v12.10.1 to fit the data. For
Mergerpn encircled energy, it was necessary to use a small re-Time after
GW170817
we fit(days)
the data with an absorbed powerlaw model (tbabs*ztbabs*powerlaw), where tbabs is agion in •order to exclude emission from the AGN, which
is only 10 arcsec from the source. We find this does
neutral absorbing column attributed to our own Galaxy,
not introduce any large systematic e↵ect when comparfixed at 8.9 ⇥ 1020 cm 2 (Dickey & Lockman 1990), and

104
10

x10 4
x10 3

110
d
165
d
172
d
209
d
218
d
297
d

x10 2
x10

2

100

m

The light curve rises and declines, which~5 orders of magnitude in γ
looks like an off-axis emission.

10

6

• is
Fermi
acceleration
works well
The spectrum
a beautiful
single powerlaw,suggesting
synchrotron
Figure2. Top: HST/F606W
light emission.
curve of the afterglow of

GW170817 spanning ⇡ 110.5 - 584.1 days (green points; observer
frame); downwards triangles denote 3 upper limits. The upper
limit at ⇡ 584.1
d is measured
from thepowers
median-subtracted
Kinetic
energytheimage,while all other data points are measured from HOTPANTS residual
images. Alsoafterglow.
shown are a structured jet model and the range oflight curves describing the top 5% of models (black solid and dotdashed lines), and a quasi-spherical outflow model (dotted line; Wu

10-2
10-4
108

5

X-ray

1010

Opticalx10 8Flux Density ( Jy)27

A single power law: ~ν-0.6
38=> synchrotron radiation(optically thin)
37.8•

Afterglow (F606W)
Structured Jet (WM18)
Quasi-Spherical(WM18)-1

Optical AB Magnitude

26.5

GW170817 HST

log(L) (erg s )26

Fong+19, also Margutti+18
RadioSpectrum

1010

1012

1014

1016

328
d
362
d
584
d

1018

Frequency (Hz)Figure 3. Broad-band SED of the afterglow of GW170817 at nineepochs of our HST observations, spanning ⇡ 110 - 584 days; fluxes
are scaled for clarity. The HST photometry in this paper (green
Hallinan+17,
circles),
radio afterglowMargutti+17,18,(red squares; Margutti Troja+17,19,
et al. (2018); Mooley
et al. Haggard+17,(2018c); Dobie et al.
(2018); Mooley et al. (2018b); Alexander
Ruan+17,Lyman+18,Mooley+18
et al. (2018); Troja et al. (2018c), Hajela et al. in prep.), and X-ray

Imaging
theafterglow
with
VLBI
Breaking the degeneracy: VLBITwo observations with the HSA(75 d and 230 d post-merger)

Superluminal Jet in GW170817
VLBI resolve the motion of the radio source Mooley…KH (2018)
1, The source moved
2.7 mas in 155 day.
Day 240

Day 75=> 2.7 mas ~ 0.5 pc (at 40Mpc)

app = 4.1 ± 0.4 at 41 M

2, The source size isunresolved.
=> the emission region
does not extend much.

Very strong evidence for a jet in GW170817• First time to see a superluminal motion of a “GRB” jet.
•Figure 1: Proper motion of the radio counterpart of GW170817. The centroid offset posi-

ions (shown by 1 errorbars) and 3 -12 contours of the radio source detected 75 d (black)Movie© Ore Gottlieb & Ehud Nakar

Sigm

2
0
-2
-4

Jet Parameters
10100
Energy (keV)

Figure 7. The 256 ms binned lightcurve of GRB 170817A in the 10–300 keV band for NaI 1, 2, and 5. The shaded regions are

1000

(a)

the different time intervals selected for spectral analysis. The inclusion of the lower energies shows the soft tail out to T0+2 s.

15 –25obody).
BGO_00 The blue bins are the forward-folded model fit to the coun

1.00
0.10

4
2
0
-2
-4

Sigma

Sigma

0.01

NAI_01
NAI_02
NAI_051.0

detector, and 2

shown in the lower subpanels.

0.1
4
2
0
-2
-4

n ≈ 10-4 – 5x10-3 cm-3
10100
Energy (keV)

(a)

1000

upper limits estimated from the model variance a

10

E ≈ 1049 – 1050 erg

100Energy (keV)(b)

Afterglowo
BGO_00
NAI_01
NAI_02
NAI_05

10.00
Rate (counts s-1 keV-1)

Figure 8. Spectral fits of the count rate spectrum for the [Left] m

GRB 170817A (promt)

≤5o

θobs
Rate (counts s-1 keV-1)

θjet

Figure 8. Spectral fits of the count rate spectrum for the [Left] main pulse (Comptonized) and [Right] softer emission (blackbody). The blue bins are the forward-folded model fit to the count rate spectrum, the data points are colored based on thedetector, and 2

upper limits estimated from the model variance are shown as downward-pointing arrows. The residuals areshown in the lower subpanels.

Fig. 10.— Eiso as a function of Ep in the burst frame for a sample of
Zhang et al. (2009). The red star is GRB 170817A. The solid line is the S
K,isowith its 2 confidence level.
togetherEinferred from VLBI
of (top
Eiso ofshort
GRBs of t
gure 5: Light curve of GRB 170817A observed byHigh
Fermiend
GBMleft).
Spectrum
Mooley+KH,18
ike of GRB 170817A (top right). Comparison of GRB 17817A with Fermi GBM GRBs: thergy and duration (bottom left) and the peak energy and isotropic equivalent energy fo• We would
strongareGRB
if we wereRBs (bottom
right).have
Theseen
top afigures
adoptedfrom on-axis.
Goldstein et al. (2017) and the b• I’ll Lu
talketabout
implications of these to CTA science in a meeting next week.
ures fromal. (2017).

GW + light curve + VLBI => H0

e Methods)15 . Figure 2 depicts the poste-

KH+19

the GW-only analysis15, 27 . The constraint+4.5
s/Mpc, to 68.1 4.3 km/s/Mpc (median and

3-4% of a systematic uncertainty due to jet modelingOutline•

A comment on Binary Black Holes: an astrophysicalimplication from new LIGO/Virgo catalogue (GWTC-2)•

Binary Neutron Star: what we have learned from multimessenger observations of GW170817•

High energy (~100 MeV) neutrinos from core collapse
supernovae and prospects with HyperKamiokande.

o the theories of stellar evolution, neutronlack holes are formed by gravitational cole stars (M8M'). Most of the gravitareleased by the collapse ( 3X 10 ergs) isutrinos. It has been argued that if neutrigravitational collapse of stars could be
ould give not only strong evidence of thee final stage of stellar evolution but also
ion on the mechanism of supernova explo-

cussion on the consistency

betweenthese two observa-

tions).
Soon after the discovery of the supernova 1987A, Bah—call, Dar, and Piran presented the expected neutrinosignals in terrestrial detectors.
In the present paper, we make an analysis of the observed burst systematically by comparing with theoretiNagakura & KH 2020
cal predictions, mainly the numerical simulation of collapse by Wilson and collaborators.
Total energy of antineutrinos
and the progenitor
e
μ
τ
mass.
In Fig. 1, the energies of eleven neutrinos detectebruary
1987, 7:35 UT the Kamioka
ed by the Kamiokande(theirexperiment
II Collaborationgroup (Kamiokande II Collaantiparticlesareasshown.
well) The
cted the neutrino burst from the supernova
energy of neutrinos E is estimated from that of electronsappeared in the Large Magellanic Cloud.
E, as E =E, +m„—mz on the assumption that all theare v, p νμ e+n,
in the Montthat events
becausethe anti-v
cross section
Castagnoli
workingacceptedof this from a supernova is
But itetisal.widely
obtainingντ andtheirspectrareaction is almost one hundred times larger than that ofclaimed that the neutrino burst was detectvery
hard
becausethey
donothave enough energy to induce charged current interactions.
the v, e scattering, and the directions of e+ (or e )uary 1987, 02:52 UT, 5 h earlier than thatformed in the detector are random except the first andande II collaboration (see Ref. 2 for a dis-)

Looking forward to a Galactic supernova

Astrophysicists’ dream is to—
get the ν ν ν spectra from a supernova.

Neutrino Time - Energy from SN 1987A (Sato & Suzuki 1987)30-~ 10 neutrinos for 2 kt & 50 kpc

I~

Mv,
10—

~ ~

e~ &~ ~~

gi ~

~

.

~ ~

~

~

Kamiokande II & LMC

~ ~

M v,

Hyper-Kamiokande & Galactic~ 105 neutrinos for 200 kt & 5 kpc
0.1

I

TIME

I

I

1

I

I

I

1

'0SeC

energies of eleven neutrinos detected by the Kamiokande II collaboration with error bars. The energies are derived
e+n process. The fine horizontal lines
b represent the mean energy
on that all the events are caused by v, p
5 a ande
flux averaged in the respective ranges (see Table II). The mean energy of neutrinos predicted by various theories are
CK, Mazurek, Cooperstein, and Kahana (Ref. 10) (v, ); H, Hillebrandt (Ref. 11) (v, ); M v„Mayle (Ref. 5) (v, );Ref. 5) (v, ); BL, Burrows and Lattimer (Ref. 15) (v, and v, ). In order to display the first event (time=0 sec) on thee scale, it was shifted artificially to the point 0.01 sec.

A question: Is this true even if we have 10 ν type neutrinos?

Looking forward to a Galactic supernova
Nagakura & KH 2020

Supernova νμ ντ and their anti-v do not have enough energy to induce charged current
interactions. Is this really true?

They need to be νμ with E>100MeV and ντ with E> 2GeV
6

Nakazato et al.

Nakazato + 15

Supernova νμ ντ have quasi-thermal spectra withT ~ 5MeV.
•ντ is impossible to induce CC interactions.

•

νμ also seems to be impossible with a quasithermal spectrum.

Fig. 4.— Neutrino number spectra of supernova with 30M! , Z = 0.02 and shock revival times of trevive = 100 ms (dotted), 200 ms(solid) and 300 ms (dashed). The left, central and right panels correspond to νe , ν̄e and νx (= νµ = ν̄µ = ντ = ν̄τ ), respectively.

result of mass accretion. After the shock revival, the acprogenitor models are adopted in total. As shown incretion stops and the mean energy of emitted neutrinos
Equation (1), their neutrino spectra, dN (M, Z, Eν! )/dEν! ,Now
the question
we ask is “Are νμ and ντ always
quasi-thermal
duringdecreases,
which is called thecoolingcore
phase.collapse?
To estimate ”
are integrated
in SRNs.
the uncertainty of the explosion mechanism, we calculate
the spectrum of SRNs for models with different values of
4.1. Supernova Explosion and Shock Revival Timetrevive (Nakazato 2013). Note that the shock revival time
The physical factors that cause the CCSN explosion

“Non-thermal” vx in core collapse
Nagakura & KH 2020

When the supernova shock is ~ 100km, the collisional mean free path of vx is O(100km).
In such a situation, some of them cross the shock multiple times and then escape (firstorder Fermi Acceleration), which has been known since Kanzas & Ellison 81.
Neutrino
shock
acceleration
in
CCSN
3
=> An observer can see accelerated neutrinos.

νμ (<100MeV)ντ (<200 MeV)
Nucleon scattering

Coherent scattering~0.2c~0.02cNucleonEnergy Sphere(~30 km)Shock(~100 km)

Nucleuse.g., Fe

This acceleration occurs as longas the shock’s optical depth ~ 1.

Note that ve do not accelerate because
CC interactions destroy them.

ation is very
Fig. 1), sinceeast up to ⇠upper limit ofocesses. Thishe energy loss
k acceleration

fractions of neutrinos escape from the post-shock flows after repeating the same process (see Eq. 3 for the condition) during which
neutrinos gain the energy from the shock wave and create the nonthermal spectrum. The reachable maximum energy is ⇠ 100 MeV
and ⇠ 200 MeV for ⌫µ , ⌫⌧ , respectively. See the text for more
Nagakura & KH 2020
details.

“Non-thermal” vx in core collapse(2)e, i.e., nucleream respeceV, the shocke↵ect in thee between up2c, we obtain
mass of tau
ciple accelerthe ⌫µ accelce the ⌫µ ab+n ! p+µ

Are νμ and ντ always quasi-thermal during core collapse? No.

νe(Quasi-thermal)νe

At CCSN sourceν τ  ν μ (and their anti-partners)
νμ100MeV

Energyντ200MeVno pair annihiremsstrahlung• Muthoseand tau have a significant non-thermal tail ~ E > 50MeV.
However,• The degeneracy
Figure2. Schematicof expected
spectrafor all flavorskickof in at 100 MeV.
betweenmu andpicture
tau is broken
becauseCC interactions.

“Non-thermal” vx in core collapse
Nagakura & KH 2020

6

H. Nagakura andK. R.Hotokezaka
18
Buras et al.: Two-dimensional hydrodynamic core-collapse supernova simulations with spectral neutrino transport

Case 1. Early post bounce ~ 10 - 30 ms

the early post bounce phase in Sec. 4.2. The similar estimation but for the later phase in the case with failed CCSN ispresented in Sec. 4.3. Finally, we discuss a possibility of muon
productions in these detectors in Sec. 4.4.

4.1 Basic assumptions
4.1.1 Analytic expression of neutrino spectrumaFigure 5. Emergent spectra of ⌫µ and ⌫⌧ at 10–30, ms after thebounce, obtained by our Monte Carlo simulations. For comparison,
we display a ⌫µ spectrum at 40–70, ms after the bounce, a thermal
spectrum of zero chemical potential with T = 5.2 MeV, and afunction given by Eq. 14. with Q = 10 MeV.

b

As described in Sec. 2, the neutrino shock acceleration does
not occur for ⌫e and ⌫¯e , indicating that their spectrum remains quasi-thermal feature (see also Fig. 2). We also notethat their average energy tends to be smaller than that of
the heavyd leptonic neutrinos, which indicates that the exponential decline of the thermal spectrum in the high energy
region is steeper than that of heavy leptonic neutrinos. For
these reasons, ⌫e and ⌫¯e in the high energy region (> 50 MeV)
would be subdominant; hence we ignore their contributionsin this estimation.
In this study, we do not directly use the raw data of theemergent neutrino spectrum computed by Monte Carlo simulations but rather fit them by an analytic formula to facilitate
intuitive understanding of the event count. In the expression, the number spectrum of neutrinos at the CCSN source(dN/dE [MeV 1 ]) is assumed to be the sum of thermal- andnon-thermal component,
✓
◆
✓
◆
dN
dN
dN
=
+
,
(12)
dE
dE th
dE nt

e

“Non-thermal” vx in core collapse
Nagakura & KH 2020

6

H. Nagakura andK. R.Hotokezaka
18
Buras et al.: Two-dimensional hydrodynamic core-collapse supernova simulations with spectral neutrino transport

Case 1. Early post bounce ~ 10 - 30 ms

Buras+06

the early post bounce phase in Sec. 4.2. The similar estimation but for the later phase in the case with failed CCSN ispresented in Sec. 4.3. Finally, we discuss a possibility of muon
productions in these detectors in Sec. 4.4.

4.1 Basic assumptions
4.1.1 Analytic expression of neutrino spectrum
As described in Sec. 2, the neutrino shock acceleration does
not occur for ⌫e and ⌫¯e , indicating that their spectrum remains quasi-thermal feature (see also Fig. 2). We also notethat their average energy tends to be smaller than that of
the heavy
leptonic neutrinos, which indicates that the expoad6 H. Nagakura andK.
Hotokezaka
18
R. Buras et al.: Two-dimensionalhydrodynamic core-collapse
supernovaof
simulations
withmsspectralneutrino transportnential decline
thermalspectrumin the high energy
2. LateSupernova~the100
Figure 5. Emergent spectra of ⌫µ andCase⌫⌧ at 10–30,
ms afterfailedtheBuras+06
region is steeper than that of heavy leptonic
neutrinos. For
bounce, obtained by our Monte Carlo simulations. For comparison,
the early post bounce phase in Sec. 4.2. The similar estimathese reasons, ⌫e and ⌫¯e in the high energy region (> 50 MeV)
we display a ⌫µ spectrum at 40–70, ms after the bounce, a thermaltion but for the later phase in the case with failed CCSN is
spectrum of zero chemical potential with T = 5.2 MeV, and awould be subdominant; hence we ignore their contributions
presented in Sec. 4.3. Finally, we discuss a possibility of muonfunction given by Eq. 14. with Q = 10 MeV.
in this estimation.
productions in these detectors in Sec. 4.4.
In this study, we do not directly use the raw data of theemergent neutrino spectrum computed by Monte Carlo simulations but rather fit them by an analytic formula to facilitate
4.1 Basic understanding
assumptions of the event count. In the expresintuitive
sion,
numberspectrumofofneutrino
neutrinos
at the CCSN source
4.1.1 theAnalytic
expression
spectrum1(dN/dE [MeV ]) is assumed to be the sum of thermal- andAs describedcomponent,
in Sec. 2, the neutrino shock acceleration does
non-thermal
not occur for ⌫e and✓ ⌫¯e , ◆indicating✓ that◆ their spectrum redN
dN
dN
mains quasi-thermal= feature (see+ also Fig.
, 2). We also note
(12)
dE
dE
dE
that their
to be nt
smaller than that ofthbe average energy tends

Cumulati

0.1

Observability: charged current interactions
Nagakura & KH 2020
occur for vμ?
0.01

0.001

120

16

Early post bounce

120

180

Failed SN late phase

H. Nagakura and K. Hotokezaka

Figure 16. Same as Fig. 12 but only for charged current reaction

100

100

10

L52 = 2

∆T=10ms

T=5MeV Q=10MeV
10 kpc(Early phase)

Cumulative Event Counts

NORMAL
Cumulative Event Counts

140
160
Energy [MeV]

Muon-Creation
1

HyperK

0.1

0.01

0.001
120

140
160
Energy [MeV]

180

120

SK
HK
DUNE
10 JUNO(Thermal)

L52 = 2

∆T=100ms

10 kpc(failed CCSN)NORMAL InV T=6MeV Q=10MeV(

Muon-Creation
1

HyperK
0.1

0.01

0.001
140

120160

Energy [MeV]

140 180
160
Energy [MeV]

180

120

Figure 17. Same as Fig. 15 but only for charged current reaction
Figure 16. Same as Fig. 12 but only for charged current reaction channels with muons.

100

L52 = 2 is
∆T=100ms
Muon production
not so promising but it can occur.

ounts

NORMAL10

trino shock acceleration occurs in the late post bounce phasefor failed CCSN,
SK in which the PNS is enveloped by a stalledshock wave locatedat ⇠ 100 km
with high mass accretion
HK
InV
DUNE
rate. Motivated by these considerations, we perform a comJUNO
prehensive
study of the neutrino shock acceleration from the(Thermal)production mechanism to their detectability by representative terrestrial neutrino observatories. The main conclusions

T=6MeV Q=10MeV
10 kpc(failed CCSN)

Muon-Creation

meanwhile ⌫µ (an⌫⌧ but sharp cut
MeV). The spect
thermal shape, si
their charged cur(ii) The observ
the condition of

0.01
0.001

SK
HK
DUNE
JUNO
100
(Thermal)Observability: ve at higher energies
60

80

100
120
Energy [MeV]

140

60

80

0.001

120

140

Energy [MeV]

60

80

100
120
Energy [MeV]

140

60

Figure 11. Energy spectra of event counts for neutrinos in the early post bounce phase. The color denotes the detector:Nagakura
SK (red), HK & KH 2020
(blue), DUNE (brown), and JUNO (light green). The event counts are integrated charged current reaction channels of each detector as
described in Sec. 4.1.3. Left and right panels show the results for the normal mass hierarchy and inverted one, respectively. For comparison,
we also display the energy spectra of events counts for thermal neutrinos as dashed lines.
Figure 14. Same as Fig. 11 but for the late post bounce

Early post bounce

Failed SN late phase

100

1000

10

∆T=10ms

NORMAL

T=5MeV Q=10MeV
SK
10 kpcHK(Early phase)
DUNE
JUNO
(Thermal)1

HyperK

0.1

0.01

60

80

100
120
Energy [MeV]

140

60

∆T=100ms

10 kpc(failed CCSN)

HyperK

10

1

0.1

0.01

0.001

L52 = 2

InVT=6MeV Q=10MeV

100
Cumulative Event Counts

Cumulative Event Counts

NORMALL52 = 2

0.001
80

SK
HK
DUNE
JUNO
(Thermal)

60100 80 120 100 140
120Energy [MeV]
Energy [MeV]

140

160

60

8

Figure 12. Same as Fig. 11 but for cumulative event counts. The integration starts from 200MeV and then proceeds
negative
Figure 15.toward
Same theas Fig.
12 but for the late post bounceenergy. The neutrino energy where each line crosses with a thin black line corresponds to the expected maximum energy of neutrinos byeach •
detector. For instance, HK is capable of detecting neutrinos up to ⇠ 110 MeV.

HyperK will see neutrinos with E>80MeV in the first 50ms.
distance
for the case with normal
masshierarchyis
5 SUMMAR• This will be a clear signature that thethresholdshock
ispropagating
inthescatteringsystematically larger than that for inverted mass hierarchy.
trum shapeisanexponential
declinearoundtheenergy,imthisstudy
detect highwith
The energy s
Thiswill
is attributed
to energy
the factneutrinos
that mostof ⌫E >
at 100the CCSN
atmosphere
~
100
km.
plying that the number of injected neutrinos is sensitive to
MeV if source
the source
is atlocatedat 10 askpc.
be massstressedgarded as a q
arrivethe Earth⌫ Itinshould
invertedhierarchy•
the temperature.
reason isthat the duration
time vx that
event
countsby thermal
neutrinos
ordersmag- in the
TheseTheveothermust
originate
fromat thethesource.
most of the(seethespectrumat E > M
in Fig.are10),
whichofresults⌧eof neutrino shock acceleration is 10 times longer than that inunitude substantial
smaller than
the totaloffor
indicatingdetectors.
that
reduction⌫ neutrinos,
events in terrestrialInµin Kazanas &Summary•

BBH spin distribution points to that field binaries are their dominant progenitors.

•

The BNS merger rate is now just the one expected, ~30/Myr in the Milky Way.

•

Kilonovae are optical-nIR emission of neutron star merger ejecta. Its heating rate ~ t-1.3 (early) and t-2.8(late).

•The GW170817 light curve agrees with the r-process heating. It requires 0.05 Msun of r-process elements
produced in GW170817.

•

The estimated rate and mass of r-process elements from GW170817 are consistent with that all r-processelements are produced by mergers.

•

GRB 170817A and its afterglow point to this merger launched a relativistic jet.

•

The VLBI measurement of the superluminal motion of the jet in GW170817 provides the Lorentz factor,
total energy, and viewing angle.

•

The VLBI measurement can be used to improve the H0 measurement, ~ 68+5-5 km/s/Mpc•

Mu and tau neutrino acceleration in supernovae occurs when the shock propagates in the scattering
atmosphere. This produces high energy tail (~100 MeV) in the neutrino spectra and breaks thedegeneracy between mu and tau. Hyper-Kamiokande will be very powerful to see these signatures.

Thank you !!!

Picture after GW170817
Shocked ejecta

Tidal tailM ~ 0.01Msunv~0.1-0.3c

nu driven wind
Viscous driven
wind

Hypermassive neutron star
(HMNS)
Lifetime ~ 10 ms-1s

Jet?

Black Hole + disk

M ~ 0.05Msunv~0.05-0.1c

Outline• Introduction• Neutron Star Merger simulation• R-process Kilonova• Afterglow Jet• Origin of binary black hole mergers