Missing Links ofHigh Energy Astrophysical Phenomena
Kazumi Kashiyama（U. of Tokyo, RESCEU）

key questions of high energy astrophysics 1
How (relativistic or non-relativistic) outflows are accelerated andhow the energy is converted to thermal and non-thermal emission?
What is the central engine? à mostly compact objects (BH, NS,WD)key questions of high energy astrophysics 1I
What kind of massive star (RSG, BSG, WR) produces
what kind of compact object (NS or BH? B field, rotation, disk?)
and what kind of explosive transient (SN, GRB or else) ?
Image credit: Chandra

key questions of high energy astrophysics III
What kind of compact binaries (BBH, BNS, BHNS, …) produces
what kind of compact object (NS or BH? B field, rotation, disk?)
and what kind of explosive transient (GRB, FRB, or else) ?
How compact binaries are formed?
Image credit: Chandra

Contents
1. a missing link
2. a missing central engine
3. a missing energy

1. A missing link

The diversity of young neutron starsal. in prep.). (Right) The ‘Mouse’ ram-pressure-confined pulsar wind nebula (1.2′ across;
NASA/CXC/SAO/B.Gaensler et al. Radio: NSF/NRAO/VLA; [19]). These images provide anindication of the variety of structures possible in PWNe.

(Young) pulsars
Kaspi 101

Fig. 2. P -Ṗ diagram for 1704 objects, including 1674 RPPs (small black dots), 9 AXPs
(blue crosses), 5 SGRs (green crosses), 3 CCOs (cyan circles), 6 INSs (mageneta squares), and 7

al. in prep.). (Right) The ‘Mouse’ ram-pressure-confined pulsar wind nebula (1.2′ across;
NASA/CXC/SAO/B.Gaensler et al. Radio: NSF/NRAO/VLA; [19]). These images provide anindication of the variety of structures possible in PWNe.

Magnetars

Kaspi 101

Fig. 2. P -Ṗ diagram for 1704 objects, including 1674 RPPs (small black dots), 9 AXPs
(blue crosses), 5 SGRs (green crosses), 3 CCOs (cyan circles), 6 INSs (mageneta squares), and 7

al. in prep.). (Right) The ‘Mouse’ ram-pressure-confined pulsar wind nebula (1.2′ across;
NASA/CXC/SAO/B.Gaensler et al. Radio: NSF/NRAO/VLA; [19]). These images provide anindication of the variety of structures possible in PWNe.

Central Compact objects
Kaspi 101

Fig. 2. P -Ṗ diagram for 1704 objects, including 1674 RPPs (small black dots), 9 AXPs
(blue crosses), 5 SGRs (green crosses), 3 CCOs (cyan circles), 6 INSs (mageneta squares), and 7

The diversity of young neutron stars
The formation rates are roughly comparable~ 1/100-1000 yr.
What makes them different?



0.3
10
1013 G
10 ms
20 s<∼ Ṁcrit,repul , a bulk of the fallback matter will be directly repelled by the spindown p

CCO formation -se, it accretes on the neutron-star surface.
Fallback accretion may be crucial!
> Ṁcrit,repul and the accretion rate is so large, the surface magnetic field of thef Ṁfb ∼e.g., Torres-Forne ́ et al.16
r can be buried and the spindown power significantly decreases (Torres-Forné et al. 2

latingthe resultsof canTorres-Forné
et al. 20161 , the threshold value for burying the ma
Fallback
accretionbury the B field
if Mdot
is smallery bedescribed
as than!

"B∗ 3/2−5−1Ṁcrit,bury ∼ 10 M⊙ s.
1013 G< Ṁfb < Ṁcrit,bury , the situation will be more complicated. At firs
ase of Ṁcrit,repul ∼∼matter accretes on the neutron star; the magnetic fields and spindown energy are con
OK, then how to make pulsars?

ar surface region. As the fallback rate decreases with time, large-scale fields emerg

down power pushes back the fallback matter. In Fig. 5, we show how the consequback accretion depend on B∗ and P for Ṁfb = 10−6 M⊙ s−1 (left), 10−5 M⊙ s−1 (center⊙s−1(right).

Let us now discuss possible connection between the diversities of neutron-star formation
accretion and the observed young neutron stars. From Fig. 5, the condition Ṁ < Ṁthe value of l. This region becomes very hot and the equation of state can be approximated by that3
intoc Rthe
ultra-relativistic gas,injectthat is,energy
Ec = 4πP
c . fallback matter via e.g., neutrino emission, pulsar an

PulsarCompetitionwind
vsthefallback
accretion
betweenfallback accretion onto
and the energy injection f

Fallback matter (Regions
II andresultin aIII)variety of outcomes and lead to a diversity of compact objects an

Constructing a new sequence of self-similar solutions

Shigeyama & KK 18
lback matter can be dividedIninto
shockedfallback matter(Region II)ofandthe surrounding
thisthepaper,
we investigate
the dynamics
fallback
matter being

reted matter (Regioninjection
III). Thefrom
motionthe surrounding
matter is
not affectedsphericallyby the energy
theofcentral
object by newly
constructing
symme

position but controlled solely by the gravity from the central object. We describe the flow in Region

We consider fallback matter marginally bound by the gravitational field oby using a self-similar solution. We note that the problems treated in this paper involves two

mass Mc and a power-law energy injection rate from the central object;nstant quantities with physical dimensions GMc and Ll .

Here we assume that the spherically symmetric flow depends on time t and r only through t

l
Q̇
=
L
t
. the pressure p take the forms ofld ξ(r, t). Thus the density ρ, the velocity v, and

Dfb

1

Ll tl− 3 D(ξ)<latexit sha1_base64="30Wgu5RZ42oT6hyGkzZADzolAbA=">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</latexit>

ρ(ξ, t) = Here Ll5/3and, l are constants and t is the time measured from the(4)onset(GMc )# Although$
GMc 1/3 our solutions are described by a single non-dimensional variablev(ξ, t) =V (ξ),(5)t
r3
Ll tl−1 Pr(ξ),
p(ξ, t) =
, ξ=
(6)
⇠c
⇠s
GMc
(GMc t2 )1/3
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nctions of ξ and t. Here we have introduced non-dimensional functions D(ξ), V (ξ), and Pr(ξ).

where r is the radial coordinate and G denotes the gravitational constantgoverning equations are described as

Pulsar wind vs fallback accretion

10

-5

10

-6

10

-7

4πD ξ

1/2

fb s

"

#$

ξ

10

-2

10

-1

10

0

1

10

-4

10

-5

10

-6

10

-7

10

-1

10

-3

fb s

10

-5

10

-7

10

-9

D

fb

ξ

s

-4

10

-3

10

-2

10

-1

ξ

!

10

0

10

1

10

2

1/2
fb s

1/2

10

-3

3/2

fb s

= 4/3

account
consistently. We should also note that
4πD ξ>10a few 100 yrs, one should!
properties
at tage ∼ξ
10and magnetic fields. We will investigate these
1/2

fb s

-5

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10

"#

-7

s

-5

10

-4

10

-3

10

-2

10

-1

10

0

10

1

10

3

10

1

10

-9

2

$

5 Summary and Discussion

5

10
fb s

& 4πD ξ
3/2

4πD ξ

1/2

-1

al.4πD2016
are one dimensional. Multi-dimensio
ξ

10

"

fb s

3/2

-3

10

We should note, however, that the self-similar

#$

3

ξ
We have presented
series of self-similar soluti
101s

the energy deposition
from a central object. T
10
-1

10

= 5/3

adiabatic
index γ in10the shocked region. In ca
!π" ξ

-1

fb

fb s

10

1

s

3/2

-7

10

-2

10

ξ &D

10
10

10

s

4πD ξ

fb s

-5

10

10

-1

10

-1

ξ

1

-3

10

)

3

&ξ

4πD ξ

10

2

fb

& 4πD ξ

fb s

1/2

10

10

1

D

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10

-9

4πD ξ

10

= 7/5

10

-7

%

ξ

10

10

-5

#$ %

0

s

10

-4

!π" ξ

10

B∗ ∼ 1013 G and P ∼ 100ms, roughly coincide

&ξ

-3

10

1

fb

10

&'(

-3

& 4πD ξ

-2

3/2

10

10

-1

fb s

-1

1

4πD ξ

10

10

10

D

3/2

0

s

fb s

10

&ξ

4πD ξ

1

fb

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10

D

= 6/5

& 4πD ξfb s

1/2

comparable population
can also
be 18
naturally
Shigeyama
& KK

)'(

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-3

#$ %

!π" ξ
arbitrarily
large fallback
accretion rate by tak
10
&'(

10

-3

10

-5

10

#$ %

-5

small. On the other10 hand, in cases of γ ≤ 4/
-7

-4

10

-3

10

-2

-1

10
ξ

10

0

10

1

10

shock front r = rs :

2

c

!
rs Q̇
For γ <= 4/3, there is no solution with accretion rate larger than Ṁcrit = GM × (4πDfb ξs )crit .
c√Fig. 2. Relations between some characteristic quantities and ξc for solutions with four different values of γ: 6/5 (top left panel), 4/3 (top right panel), 7/5(bottom left panel), and 5/3 (bottom right panel). l = 0 for all the solutions.

In general,
the× maximum
accretion∗⊙ rate that can be repelled by the central energy
rs ∼ 4.2
10 cm ξs c fallback
.
(29)
√
20
s
2 1/3
∼ce12
Whenwiththe fallback
accretion
theaspositionof Since
the shockwill be ras ≈ξs (GMaccretion,is km.
described
Mc , rs , (4πD
Eq. (27).
we consider
fallback∗ tfb )
fb ξstarts,s ) and Q̇
he critical situation, the inner contact radius should be set as the neutron star radius, i.e., ξc,crit ≈a neutron star, we take Mc = M∗ ∼ 1.4 M⊙ . Also we set the radius of the neutron star as!
"2/3(GM∗ tfb 2 )1/3 , or
& KKt 18
2 1/3t
∼ 12 km. When the fallback
accretion
starts, the position of the shock willShigeyama
be rs ≈ ξs (GM,fb∗ fb )
9rs ∼ 4.2 × 10 cm ξs!
(29)"−2/3.
t20
In ξthe criticalsituation
fb sof NS with fallback,
−4
.
(30)
c,crit ∼ 2.8 × 10
20s"
he critical situation, the inner!contact
tfb 2/3radius should be set as the neutron star radius, i.e., ξc,crit ≈9
r2s ∼× 10 cm ξ.ß the shock surface at the fb time(29)
1/34.2temperaturein(GM, orthe shockeds region∗ tfb )
20 s is typically high and the radiation pressure dominates so that!to and"−2/3
adiabatic
index
will thebe closeslightlylargerthan
4/3.asCombining
Eq.(30),
2this
4 fact3 withi.e.,
2 wehesituation,
inner
contact
radiusshouldbeset
theneutronstar
radius,ξ
tfb starc,critioncriticalrate froma rotating
neutroncanbedescribed
as
Q̇
≈
(µ
Ω
/c
)
×
(R
/R
)
,≈
or
−4
crit
lc
∗
√
ξ
∼
2.8
×
10
.
(30)
ß
the
contact
surface=theNSsurfacec,critmate) 2fromas! fb ξs"! Fig. "2−220 s(4πD(GM∗thetfb 2critical)1/3 , orvalues of ξs and
B
P
∗
45
−1
#
Q̇
∼
2.7
×
10
erg
s
, radiation pressure dominates so that (3
crit
"−2/3
13 G
temperature in the shocked!regionistypically
high
and
the
10
10
ms
ξs,crit ∼ 0.3 and −4(4πDtfb(31)
fb ξs )crit ∼ 10.ß γ >~ 4/3
2 4 32
ξ2c,crit
∼
2.8
×
10
.
(30)
ction
rate
from
a
rotating
neutronstar
canbedescribed
as
Q̇
≈
(µ
Ω
/c
)
×
(R
/R
)
,
or
crit
lc Eq.(30),∗close tothanis4/3.
Combining
this fact withwe
20ands !slightly
eadiabaticµ = B∗ index
R∗ 3 iswill
thebemagnetic
moment,Ω"larger=!2π/Ptheangularfrequency,
andR
lc = c/Ω is th"2> 10 s. Finally, the maximum energyB∗√ξvalid
P Fig.
eafterthe
wecriticalassumevalues
l = 0ofwhichistypicallyfort ∼−22tfbas∼
45ξ and−1(4πDmate)
froms region
fb energystemperature
is typicallyhighand theradiationpressure
dominates
so thata splitmonopolelikespindown
Q̇
∼ the
2.7
×note
10 thaterg
sthe above,ßrate(32)cylinderradius.
Weshocked
injectionis different
fromtheclassical
dipo
crit in
13
10
G
10
ms
10
#
2 4 to3and slightly larger than 4/3. Combining
2 Parfrey
et withal.
16Eq.(30),adiabatic
index
will
be
close
this fact
we
2
3
own
rate,
Q̇
≈
(µ
Ω
/c
).
The
additional
factor(R/R)
represents
theenhancementlc∗
re µ =
Bξ∗s,crit
Rdipoleisthemagnetic
moment,
Ω
=
2π/P
is
the
angularfrequency,
and
R
the
∼
0.3
and
(4πD
ξ
)
∼
10.
∗
lc = c/Ω is(31)fb
s crit√mate
thecriticalof ξthatandξsBeloborodov) fromFig. 2 rate
assSpitkovsky,fb &own
luminosity(Parfrey,2016);
the magnetic
are maximal
t cylinder
radius.values
We notethe(4πDabove
energyinjectionis different
from thefields
classical
dipole>eafter we assume l = 0 2whichis#typically valid for t ∼ tfb ∼ 10 s.
Finally, the maximum energy
4 3
2
down
rate,
Q̇
≈
(µ
Ω
/c
).
The
additional
factor(R/R)
represents
thetheenhancementoflike a split
due tofb the
From Eqs.lc(29-32),
we obtain
critical accretio
dipole∗ξs,critmonopole,∼ 0.3 and (4πDξs )accretion.
(31)
crit ∼ 10.
10 if Mdot is smaller than
Fallback
accretion
canberepelled
the magnetic fields are maximallysdown luminosity (Parfrey, Spitkovsky, & Beloborodov 2016);>
eafter we assume l = 0 which is typically valid for t√∼ tfb ∼ !10 s. Finally,
maximumenergy"2 ! the"−2 ! accretion"2/3
n, like a split monopole, due−5to the accretion.
From
Eqs.
(29-32),
we
obtainthecritical
ξ
(4πD
ξ
)
B
P
t
s,critfb
s crit∗fb
Ṁcrit,repul ∼ 8 × 10 M⊙ s−1
.(3
10
13
as
0.3
10
10 G
10 ms
20 s
√
!
"2 !
"−2 !
"2/3
< Ṁcrit,repul , a bulk of −5(4πDwill)critdirectlyB∗ repelled
P by the spindowntfb
fb ξsbethe fallback
matterpowe−1 ξs,critb

Pulsar wind vs fallback accretion

Fallback accretion onto NS
Should depend on the progenitor inner structure e.g., Ugliano et al. 12; Ertl et al.16

The fallback typically starts when the neutrino luminosity significantly decreases, i.e., tfb ∼ 10 safter the core bounce, and the accretion rate subsequently decreases as ∝ (t/tfb)−5/3 for t ∼> tfb .
The total mass of the fallback matter increases for more massive stars, ranging fromMfb ∼ 10−(2-4) M⊙ .
Correspondingly, the peak mass accretion rate is estimated to be ∼ 10−(3-5) M⊙ s−1.

The fate of newborn neutron stars
with fallback accretion

more abundant,less massive progenitor

less abundant,
more massive progenitor

The fate of newborn neutron stars
with fallback accretion
Low field magnetar?

Crab?

1987A?

PSR?
CCO?
magnetar?

High field magnetar?

Cas A?

The fate of newborn neutron stars
with fallback accretion

typical ranges of parameters for newborn NSs ~ the intersection of the boundaries
à The fact that the three classes have a comparable population can be naturally explained?

2. A Missing Central Engine

Fast radio bursts• τ ~1-10 ms• Sν ~ 0.1-1 Jy• ν ~ GHz• event rate ~ 103 day-1 sky-1• δtν ~ ν -2
• DM ~ 500-1000 cm-3 pc• ~50 events so far with Parkes,
Arecibo, and GBT, UTMOST,
ASKAP, (and CHIME?) but notwith LOFAR• One is repeating (FRB121102)!
Figure 2: Frequency evolution and integrated pulse shape of the radio burst. The survey data,collected on 2001 August 24, are shown here as a two-dimensional ‘waterfall plot’ of intensity
as aal.
functionradio frequency versus
time. The dispersion
clearly seen
as a2014;
quadratic sweep
Lorimer et al. 2007; Keane et al. 2012; Thornton et2013;ofBurke-Spolaor& Bannister
2014;isSpitler
et al.
across the frequency band, with broadening towards lower frequencies. From a measurement ofRavi, Shannon & Jameson 2015; Petroff et al. 2015;
Masui et al. 2015; Champion et al. 2015 …
the pulse delay across the receiver band using standard pulsar timing techniques, we determine
the DM to be 375±1 cm−3 pc. The two white lines separated by 15 ms that bound the pulse show
the expected behavior for the cold-plasma dispersion law assuming a DM of 375 cm−3 pc. The

FRB 121102

LETTERRESEARCH
Spitler+16,
Chatterjee+17, Marcote+17,Tendulkar+17, …
a
33° 12′

Declination

It repeats! e.g., ~10 times in ~3 hrs

The host is identified!

33° 9 ′

33° 6′

a dwarf star-forming gal.
@ z = 0.19
5 h 32 min 12 sb250

5 h 31 min 36 s

Observed
Average
FRB detection

225

Flux density (μJy)

5 h 32 min 0 s
5 h 31 min 48 s
Right ascension

200
175
150
125

A persistent radio counterpart!
5

10

15

20

25

30 35
125 130
MJD – 57,500 (days)

135 140145 150

The Persistent Radio Counterpart
RESEARCH LETTER
Spitler+16, Chatterjee+17, Marcote+17,Tendulkar+17, …
Energy (eV)Marcote et al.

The Astrophysical Journal Letters, 834:L8 (9pp), 2017 January 10

a

10 –5

10 –4

10 –3

10 –2

10 –1

10 0

10 1

10 2

10 3

10 4

33° 12′
Henize 2-10
Radio-loud AGN
Crab nebula

10 –13

UKIDSSGLIMPSE
Gemini i-band
Keck R-band

33° 9 ′QFQ (erg cm–2 s–1)

Declination10 –14

Gemini r-band

XMM-Newton
and Chandra

10 –15ALMA
10 –1633° 6′

VLA
10 –17
Radio counterpart
Optical counterpart
10 –185 h 32 min 12 s

5 h 32 min 0 s

5 h 31 min 48 s

5 h 31 min 36 s

109

1010

1011

1012

1013

1014

1015

1016

10 17

10 18

Frequency (Hz)

Flux density (μJy)

Figure 1. EVN image of the persistent Rightsourceascension
at 1.7 GHz (white contours)
Figure 3 | Broadband spectral energy distribution of the counterpart.
together with the localization of the strongest burst (red cross), the other three250bursts (gray crosses), and the position obtained after averaging allbobserved
Detections of the persistent radio source (blue circles), the optical
Observed
four bursts detectedAverageon 2016 September 20 (black cross). Contours start at a 2σcounterpart (red and orange squares) and 5σ upper limits at various
FRB
detection
noise level
andincrease by factors of 21 2 . Dashed contours represent
225 of 10μJyfrequency bands (arrows) are shown; see Methods for details. Spectral
negative levels. The color scale shows the image at 5.0 GHz from 2016energy distributions of other radio point sources are scaled to match theSeptember 21. The synthesized beam at 5.0 GHz is represented by the gray200
radio flux density at 10 GHz and overlaid for comparison: low-luminosity
ellipse at the bottom left of the ﬁgure and for 1.7 GHz at the bottom right. The
lengths of the crosses represent the 1σ uncertainty in each direction. Crosses for
AGN in Henize 2-10, a star-forming dwarf galaxy28 placed at 25 Mpc
Kokubo+17
each individual
burst
reﬂectonly
thestatisticalerrorsderived
from
their
S/N
175
(blue); radio-loud AGN QSO 2128− 12329 scaled by 10− 4.3 to simulate aand the beam size. The size of the cross for the mean position is determined
Figure 2.lower-luminosity
Top: dynamic spectrum
the placedstrongestatburst
on 2016
AGNofand3 Gpcdetected(yellow);
and the Crab nebula30
from the spread of the individual burst locations, weighted by ξ (see the text),
September 20 (burst #2 in Table 1) from Arecibo autocorrelations, showing

The source is localized within ~ 0.7 pc (!!)& associated with a star forming region …

The spectrum is compatible with
the Crab pulsar-wind nebula …

An FRB in a bottle?

Or a bubble?
= pulsar wind nebula& supernova remnant

A Very Young NS in a Bubble~ a few months after the explosion~ 1-100 yr after the explosion

The PWN emission is absorbed
and thermalized in the supernova ejecta,
powering a luminous supernova.

The non-thermal pulsar wind nebula (PWN) emission in
the radio bands starts to escape the supernova ejecta.
à repeating FRBs and the persistent radio counterpart?

pulsar wind nebulasupernova ejecta

A Very Young NS in a Bubble~ a few months after the explosion~ 1-100 yr after the explosion

The PWN emission is absorbed
and thermalized in the supernova ejecta,
powering a luminous supernova.

The non-thermal pulsar wind nebula (PWN) emission in
the radio bands starts to escape the supernova ejecta.
à repeating FRBs and the persistent radio counterpart?

pulsar wind nebulasupernova ejecta

Very young, but not too young
Energy (eV)
10 –5

10 –4

10 –3

10 –2

Henize 2-10
Radio-loud AGN
Crab nebula

10 –13

10 –1

10 0

10 1

10 2

10 4

UKIDSSGLIMPSE
Gemini i-band
Keck R-band

10 –14

QFQ (erg cm–2 s–1)

10 3

Gemini r-band

XMM-Newton
and Chandra

The flux is much higher thanthe Crabà Much younger and powerfulthan the Crab pulsar.

10 –15ALMA
10 –16In the early stage, radio waves
cannot escape the nebula ...

VLA
10 –17
Radio counterpart
Optical counterpart
10 –18

109

1010

1011

1012

1013

1014

1015

1016

10 17

10 18

Frequency (Hz)Figure 3 | Broadband spectral energy distribution of the counterpart.
Detections of the persistent radio source (blue circles), the opticalcounterpart (red and orange squares) and 5σ upper limits at various
frequency bands (arrows) are shown; see Methods for details. Spectralà The NS is sufficiently old and/or
the SN ejecta mass is small.

As we will show in a separate paper, the bright persisent radio emission from the pulsar wind nebula is also
aturally explained within this scenario.
The population of repeating FRBs is consistent withs follows.
The rate
repeatingrate
FRBs
inferredhedouble neutron
starof formation
withinuncer4
1 FRB 1sources canainties.
Theis populationof
repeatinge survey
5.1+17.8
·
10
sky
day
(Scholz
4.8
e16).
estimatedbasedon thatthe Arecibo’s
survey
area
Given
the
11 bursts
arefrom
detected• The
rate
of fact
repeating
FRBs
inferredtheandsurveyme. The rate of repeating FRBs inferred from the suray
andFRBs
by 2016).
the(Scholz+16)
ey is⇠ repeating
1.4 · 10 5 sq.
deg 1ares 1 detectable(Scholz et al.
Given
the fact
that 11 bursts
are⇠detected
in 0.6day
telescopeatdistances
out
to
1
Gpc,
the
for• which
cantimebe translated
to the formation
rate
as
nd
the
surveyfor
eachpointing
is
233
s,
a
corrate of repeating FRB sources is roughly estiection factor of1 ⇠ 0.051 should
be taken into1 account.
1
ohereforebe ⇠ 60f
Gpc rate
yr of(T/10 yr) . HereFRB sources is estimatedthe bformation
FRB
1s⇠
700fb 1lifetime
Gpc 1 yr(⌧FRB /10 yr), where⌧FRB
is aatypicalof 1repeating
FRB
objectsandypicallifetime
fb Tis~a0.1
beamingfacis a beamingof& them.
⇠
10yr,
for τFRBoffactor~FRB
10 yrobjects
fb ~ and
1 Ifà
%
of
CCSNe
or of FRBs. If fb ⇠ 1 and ⌧FRB = 30FRB
yr, the rate is
forwithτFRB ~
100ofyrthe&formation
fbformation~ 1 à rate~of0.01% doubleof CCSNeis compatible
withtheof binaryompatible
that
close
(though
uncertainty
is forhuge)
stars (see,e.g.,Kalogeraal.
2004;
et al.
eutronstars
withthe(Kim
et al.et2015
the Kimgalactic
douletheneutronstars and
Wanderman
Piran2015 forandshort
Galactic
doubleneutron&star
systems
GRBs). The rate is also consistent with that of superman & Piran2015forfor⌧ shortGRBs).
uminoussupernovae=
100
yr. These lifetimes
FRB

Formation rate

21102
ivistic
PWN
an be
(11)

Constraints on NS & SN
KK & Murase 17

ons inmax istimate(12)

ctor iny lossü PWNe with Mej ~ a few Msun, P0 ~ ms, Bdip ~ 1013 G, & tage ~ a few 10-100 yrs
e also can explain the repeating FRB.
t. Theü If tage ~ 100 yrs, the formation rate is ~ 0.01 % of core-collapse SNe.
an be

A Very Young NS in a Bubble~ a few months after the explosion~ 1-100 yr after the explosion

The PWN emission is absorbed
and thermalized in the supernova ejecta,
powering a luminous supernova.

The non-thermal pulsar wind nebula (PWN) emission in
the radio bands starts to escape the supernova ejecta.
à repeating FRBs and the persistent radio counterpart?

pulsar wind nebulasupernova ejecta

Diversity of pulsar-driven supernovae“Height” Multi-Messenger Tests for Fast-Spinning“Width”
Newborn
NSs in SE SNe
Kashiyama et al.

-19

-19
-18

-18
-17

-17
-16

-16
-15

-15
-20
-40

P0 = 1 ms, Bp = 1×1013 G, Mej = 10 Msun
P0 = 1.5 ms, Bp = 3.5×1013 G, Mej = 7.5 Msun
P0 = 2 ms, Bp = 1×1014 G, Mej = 5 Msun
P0 = 2 ms, Bp = 3×1014 G, Mej = 5 Msun

-10
-20

0
10
20
30
20
40
60
80
100
Days since Maximum
Days since Maximum
Fig. 3.— Comparison of observed SN Ibc and BL-SN Ibc (greyFig.
2.—light curvesof pulsar-driven
SN andmodelswith mslines)and SN
theoreticallight curves.
The red solid
dotted-dashedrotationand diﬀerent
magnetic-fieldstrength.
The upper
lower
lines show
wind powered
cases, whereas
the bluedashedand
lines
show
0
grey56lines
show theobserved SL-SN-Ic
PTF 09cnd and BL-SNe Ic
Ni-powered
cases.
1998bw, respectively.56 Ni
ej 20

For a given P ,estimated
Lsn ≈ L Ni× (tbrightest
SNasbecomestheNi /tesc ) , orif
56

56

"!
!
Actually, e.g., Maeda et al. (2007)
proposeda"
newborn
−1
56
56 Ni
M
M
ej
Ni as a relevant
42−1
magnetarenergysource
of
a
type
Ibc
SN
Lsn ∼ 4 × 10 sd
erg s
dif
0.1
M
5
M
⊙ to show⊙ in what
2005bf. In this regard, it is important!
"−1! SN Ibc is"compatible
parameter range ordinary
KT with the
Vej
(12).
× and 9how to−1identify the−1
pulsar-driven scenario
underlying
10 cm s0.05 g cm2t⇠t

B Magnitude-20

MR [mag]

MR [mag]-21 -20In this work, we numerically calculate light curves of0
SNe driven
by fast-spinning strongly-magnetized newborn NSs embedded in SE progenitors. Details of themodel description are given in Appendix. We assume
that the energy injection is caused by spherical winds
rather 0.5
than jets, and both 56 Ni decay and magnetized
wind are taken into account as energy sources. The
thermalization of the non-thermal emission is approximately taken into account, and the optical SN emission and early non-thermal nebular emission are ob1
tained consistently.
The eﬀect of GW spin-down is incorporated in a simple parametric form. The present
14
Esn = 1×10
erg, Mej = 2 Met
ms, Bp = 5×10but
G with sevmodel is basedon 51Murase
sun, Pal.
0 = 10 (2015)51
13
Esn = 1×10 erg, Mej = 2 Msun, P0 = 10 ms, Bp = 5×10 G 56
eral refinements,
e.g.,
including the eﬀect of
Ni decay.
51
14
Esn = 1×10 erg, Mej = 5 Msun, P0 = 10 ms, Bp = 5×10 G
The simple
model
allowsus
toexploreawideparame1.5
ter range;
1 − 30 ms,-20the initial-10 spin of
0 Pi = 10
20 poloidal30
13−15magnetic field of Bdip =Days
10since Maximum
G, SN ejecta mass of56= 0.05−1.0,
SN exMFig.
Ni mass of M56 Ni SN
ej = 1−10 M⊙ ,4.— SN light curve of pulsar-driven
models around the51−52
plosionenergyof Esnwith= 10erg, andbody
opacpeak. The
grey lines
arrows indicate
the greyobserved
range
of
−1
−2
SNeKIbc
and
BL-SN
Ibc.
ity
=
0.05
−
0.2
g
cm
.
Note
that
K
∼
0.1
and
For a given P0 , TT
0.2 g−1 cm−2 corresponds to electron scattering for singly
SN helium,
becomesslowerand can be
ionized and fully ionizedrespectively,
smaller for e.g.,
partially
Mej = 2aMsun, MNi = 0.05ionized
Msun, Εsn C= 1 ×or
1051O-dominated
erg
for
a
smaller
B15
p
23
ejecta.10.0=1
In Fig.
2, we , 15show
some light curve 22exSN Ibc
MR
amples of the millisecond-pulsar-driven
SN +16model.
KK
21
The thicker red lines correspond to larger magnetic
fields. The grey lines indicate the observed SL-SNEK = 1052 erg-22 -2152

Esn = 3×10 erg, Mej = 8 Msun, MNi = 0.8 Msun
51
Esn = 6×10 erg, Mej = 4 Msun, MNi = 0.4 Msun
51
Esn = 2×10 erg, Mej = 2 Msun, MNi = 0.2 Msun
51
Esn = 1.5×10 erg, Mej = 2 Msun, MNi = 0.05 Msun
51
14
Esn = 1×10 erg, Mej = 5 Msun, P0 = 2 ms, Bp = 3×10 G
51
14
Esn = 1×10 erg, Mej = 2 Msun, P0 = 20 ms, Bp = 7.5×10 G

3

MR-MR,max [mag]-23 -22]4

Multi-Messenger Tests for Fast-Spinning Newborn N

1015

Mej = 5 Msun, MNi = 0.1 Msun, Esn = 3 × 1051 erg
23=

BL-SN Ibc
.0

SN Ibc12221- Peak R-Band AB Magnitude

Superluminous Supernovae – decline
FRBs?
phase, or

ter the peaks. M
KK+16
late-time optical
Murase, KK, Meszaros 16

4.
See also
Metzger+16, Margalit+17

erg

EK = 10 52

Bp [G]

Because of add
SN model, optica
20
distinguish the m1014
19els. To break pa
approaches are t
18
nario for SE-SN
SL-SN Ic
17
and SLSN Ic, an
Ni poweredtion to SE-SNe.
161013
sar engine is the1
10
in x rays (e.g.,
P0 [ms]Murase et al. 201
MuraseetGal. 20
13
ü
Pulsar-driven
superluminous
SNe:
M
~
a
few
M
,
P
~
ms,
B
~
10
ej
sun
0
dip
Fig. 6.— Same as Fig. 5, but with Mej = 5 M⊙ , M56 Ni =nism of magneti
51
àconsistent
with
the
young
NS
model
for
FRB
121102
0.1 M⊙ , and Esn = 3 × 10 erg. Parameter regions broadly conlikely outcome isistent with
Ibc, BL-SN~ Ibc,
Ic are SNeü theTheobserved
event rateSN
is consistent;
0.01and% ofSL-SN
core collapse
trons, which trig
indicated.
ü The host gal. type is also consistent
synchrotron emisthat proto-NSs with relatively weak poloidal fields can5
,1RM

M R, 15 = 0.3

A Very Young NS in a Bubble~ a few months after the explosion~ 1-100 yr after the explosion

The PWN emission is absorbed
and thermalized in the supernova ejecta,
powering a luminous supernova.

The non-thermal pulsar wind nebula (PWN) emission in
the radio bands starts to escape the supernova ejecta.
à repeating FRBs and the persistent radio counterpart?

pulsar wind nebulasupernova ejecta

Radio PWNe of SLSN remnants
6

Omand et al.

1GHz

Omand,KK,Murase 17

0. fit SLSN light curves with the pulsar driven model
1. calculate the early PWN emission

100 GHzü Detectable with ALMA from ~ 1 Gpc<~ a few yrs after the explosion
ü Detectable with VLA from ~ 1 Gpc~ few 10 yrs after the explosion, andconsistent with FRB121102

Still missing …
• The observed upper limits on the radio PWNe
by ALMA and NOEMA

pre
lim
ina
r

y

by VLA (including FRB search)

pre
lim
ina
r

y

Omand et al. in prep
Law et al. in prep

Powerful survey facilities online
UTMOSTan upgrade of the Molonglo Observatory
Synthesis Telescope (MOST)CHIMEThe Canadian Hydrogen Intensity Mapping Experiment

ASKAP
The Australian Square Kilometre Array Pathfinder

HIRAX
The Hydrogen Intensity and Real-time Analysis eXperiment3. A Missing Energy

Ia SNe?

Double WD merger

Not found~ 1 yr -1gal-1
Debris expansion~ 1050erg/100yr~ 10 % of CCSN

Fast-spinning WD@ mass shedding limit?

à Prot = 1-10 sec

Cosmic ray e factories?

Kashiyama, Ioka, & Kawanaka 11

Origin of FRBs?

Kashiyama, Ioka & Mészáros 13

Let’s search for fast spinning single WDs!

The Target Signal
e.g., magnetized (3x108 G)WD G111−49

See also

P = 6.68 hr
Amp. ~ 5%• DECam minute cadence survey Belardi+ (2016)• Kepler/K2 observation of pulsating WD Hermes+ (2017)

But there is no sub-mimute cadence survey so far …

How many fssWDs?
WD birth rate:

ṄWD ⇠ 1 yr 1 gal 1 Badenes & Maoz (2012)

Double WD merger rate: Ṅmerger ⇠ 10 (2-3) yr 1 gal 1
à fraction of merger-origin WDs: f↵sWD ⇠ 0.1-1%
Local number density of WDs from Gaia observations

nWD = 4.49 ⇥ 10 3 pc 3

à4.49 ⇥ 10

5

pc

3

✓

Hollands+ (2018)

ffssWD
1%◆✓

nWD
4.49 ⇥ 10 3 pc◆

The Hertz Spinning Object surveywithWhy with Tomo-e?
Limiting Magnitude (SNR = 5)28
Tomo-e
DeCam
HSC
LSST
Pan-Starrs
ZTF

26
24
22
20
18
16 1
10

100

101
Time Resolution [s]102Because it’s a CMOS camera
capable of sub-minute cadence!

103& because of its wide field of view!

Prospects

Number of detected WDs per bin• Limiting magnitude: g = 19, sky coverage 10,000 deg2
• ffssWD = 0.3%
16
Bp sin ↵ = 108 G

14

Bp sin ↵ = 109 G

12
10
8
6
4
2
0
1.0

1.5

2.0
log period [s]

2.5

3.0

Pipeline Construction,
Test Observations Under Way

pre
lim
ina
r
WD
 K2 field HeSO ASF

ySummary
1. The diversity of young neutron stars is originatedfrom the fallback accretion onto the newborns?
2. The central engine of (repeating) FRBs andsuperluminous SN is the same?
3. Let us search for the fastest spinning white dwarf!