Basic reviews on
Galactic cosmic-ray origin
1
霜⽥治朗cf.) Hayakawa et al. 1958 （伝統的・標準的な話題）
Gabici et al. 2019
（宇宙線業界の最近の話題）

1, 名古屋⼤学
夏の学校2020, 8⽉26⽇, 星間現象

宇宙線（Cosmic Rays）• 広義：宇宙からやってくる「放射線」（光⼦，核⼦，
電⼦，ニュートリノ，etc…）• 狭義：「荷電粒⼦」（電⼦・陽電⼦，陽⼦，ヘリウ
ム核，etc…）→ 天体物理学では，特に光速に近い速度のものを指すことが多い．
• Victor Franz Hess博⼠(1883-1964)が1912年に発⾒し，1936年にノーベル物理学賞を受賞した．

宇宙線（Cosmic Rays）• 広義：宇宙からやってくる「放射線」（光⼦，核⼦，
電⼦，ニュートリノ，etc…）• 狭義：「荷電粒⼦」（電⼦・陽電⼦，陽⼦，ヘリウ
ム核，etc…）→ 天体物理学では，特に光速に近い速度のものを指すことが多い．
Ø
宇宙線の起源は発⾒から100年来，謎のまま．
• Victor
Franz Hess博⼠(1883-1964)が1912年に発⾒し，
Ø1936年にノーベル物理学賞を受賞した．
今⽇は60年代頃までに提唱された「標準シナリオ」を紹介する．（cf. Hayakawa, Ito, Terashima 1958;
Ginzburg & Syrovatskii 1964）Ø 最近の話題は Gabici et al. 2019 などを参照．

5th Fermi Symposium : Nagoya, Japan : 20-24 Oct, 20143

3.2. Gamma-ray Flux Predictedby P.Morrison [4]From Nishimura 2015

宇宙線（Cosmic Rays）

Six year after the perdition by Hayakawa, P
Morrison advocated the importance of Gamma ray
astrophysics in 1958. His main argument is that theastrophysics was developed in the past by observing
visible light and radio wave, but those photons were
早川幸男先⽣（1923-1992）．
descendants of the gamma rays produced from the high宇宙線の超新星起源説，およびガンマ
energy phenomena in the stars and Galaxies. In this線天⽂学の提唱者．
respect, it is important to observe directly the gamma
rays from the source to understand what are happening
天⽂学会の早川幸男基⾦（若⼿海外学in the space. Instead of diffuse gamma rays estimated
術研究援助基⾦）を設⽴された．
by Hayakawa, he focused to the point sources of gamma（天⽂辞典より）rays of the Active stars and Galaxies.
He first discussed on the processes of gamma-ray
production:
Synchrotron, Bremsstrahlung, Nuclear gamma rays, πodecay, Matter and Antimatter annihilations.
In the case of the Radio luminous colliding galaxies
of Cyg-A, he estimated the gamma ray flux by
assuming the energy source of galaxy is due to thematter – antimatter annihilation. His estimated flux of
Gamma rays of Cyg-A was
0.1-1.0/cm2sin the range of a few MeV to a few hundred MeV. Thisis several orders higher than Hayakawa’s estimate forthe flux of diffuse gamma rays.
Then Morrison proposed we could observe the
gamma rays rather easily, if we point the detectors to• 広義：宇宙からやってくる「放射線」（光⼦，核⼦，
電⼦，ニュートリノ，etc…）• 狭義：「荷電粒⼦」（電⼦・陽電⼦，陽⼦，ヘリウ
ム核，etc…）→ 天体物理学では，特に光速に近い速度のものを指すことが多い．
Figure2: S. Hayakawa
Ø
宇宙線の起源は発⾒から100年来，謎のまま．
performed comprehensive
works on Cosmic rays and high•HeVictor
Franz Hess博⼠(1883-1964)が1912年に発⾒し，energy Astrophysics, and predicted the importance of
1936年にノーベル物理学賞を受賞した．
Gamma-ray
astrophysics through  decay in 1952.
Ø
今⽇は60年代頃までに提唱された「標準シナリオ」
oを紹介する．（cf. Hayakawa, Ito, Terashima 1958;
Ginzburg & Syrovatskii 1964）

he discussed how much interstellar matter is traversedby cosmic rays during he transport from the sources tothe Earth, referring to the observed data on heavy
primaries in cosmic rays by two groups of the Bradt –Peters (1950), and Dainton-Fowler-Kent. (1951).
In his paper, he also mentioned that gamma-ray
emission from the πo mesons produced in collisions ofØ 最近の話題は Gabici et al. 2019 などを参照．

宇宙線（Cosmic Rays）
CE²ĕīĽŐĮŔĶńįĻŐThe energy spectrum of CRs

CE²ĕīĽŐĮŔĶńį
ĽŐĮŔĕOhira, Yamazaki & Terasawa (2012)­@, e-, p, He, …

dF/dE∝E-2.7

ðīĽŐĮŔĕ
-, p, He, …ýĠ100Uþ¯Ď¹è­@, e15.5

< 10

eV

> 1015.5 eV

SNROrigin？？？
ÀýĠ100Uþ¯Ď

17eVĖ®;?!
017eVĖSNR?!E>10ðīĽŐĮŔĕ

17eVĖSNR?!E>1017eVĖ®;?!
E<10¹è­@, e-, p, He, …
=T,PQ,HñB Ç, 2012

=T,PQ,HñB Ç, 2012ÀýĠ100Uþ¯Ď宇宙線（Cosmic Rays）
CE²ĕīĽŐĮŔĶńįĻŐThe energy spectrum of CRs

CE²ĕīĽŐĮŔĶńį
ĽŐĮŔĕOhira, Yamazaki & Terasawa (2012)­@, e-,Ionizationp, He, …

ðīĽŐĮŔĕ
ýĠ100Uþ¯Ď¹è­@, e-,f p, He, …

forSuper Novae
e
g
n
g
i
Tr
tio
a
ÀýĠ100Uþ¯Ď
rm
17
o
f
0 eVĖSNR?!E>1017eVĖ®;?!

ðīĽŐĮŔĕ
17eVĖSNR?!E>1017eVĖ®;?!
E<10-, p, He, …
¹è­@,eMolecular=T,PQ,HñB

CloudÇ, 2012

=T,PQ,HñB Ç, 2012ÀýĠ100Uþ¯Ď宇宙線（Cosmic
Rays）
宇宙物理学での宇宙線の例
CE²ĕīĽŐĮŔĶńįĻŐThe energy spectrum of CRs
1. ISMのエネルギー平衡 (太陽近傍)
Ohira, Yamazaki & Terasawa (2012)

CE²ĕīĽŐĮŔĶńį
CR ~ turbulence ~ B-field ~ thermal ~ 1 eV/cc
ĽŐĮŔĕ→ ガスの運動に影響を与える:­@, e-,Ionizationp, He, …
e.g. 銀河⾵の駆動源 (Breitschwerdt+ 1991)ðīĽŐĮŔĕ
2. 分⼦雲の電離源ýĠ100Uþ¯Ď¹è­@, e-,f p, He, …

forSuper Novaee
電離ガスは磁場と結合している．
g
n
g
i
Tr
tio
a
ÀýĠ100Uþ¯Ďrm
17o
磁場はプラズマ状態に依存して散逸する．
f0 eVĖSNR?!E>1017eVĖ®;?!

ðīĽŐĮŔĕ
17eVĖSNR?!E>1017eVĖ®;?!
→ 星形成中の磁束と⾓運動量の再分配(cf. Inutsuka 2012).
E<10-, p, He,
¹è­@,
e
…
Molecular=T,PQ,HñB
Ç,2012
宇宙線の起源は？

Cloud=T,PQ,HñB Ç, 2012ÀýĠ100Uþ¯Ď宇宙線の起源について①
CE²ĕīĽŐĮŔĶńįĻŐThe energy spectrum of CRs-3 ( mostly
Ohira,
Yamazaki& Terasawa(2012)
Ø εCE²ĕīĽŐĮŔĶńį
~1eVcm
reserved
in ~GeV CRs)
CR
＊銀河内の前エネルギーの推定：ĽŐĮŔĕ
55 erg (V/ 1067 cm3)
~
ε
V
~
10
­@, e ,CRp, He, …ðīĽŐĮŔĕ
Ø 銀河が¹è­@,“準定常状態(quasi-equilibrium
state)”にある
ýĠ100Uþ¯Ďe , p, He, …
とすると， これを維持するために必要な宇宙線の注⼊率が存在する．
ÀýĠ100Uþ¯Ď1717eVĖ®;?!
0 eVĖSNR?!E>10→ PCRðīĽŐĮŔĕ
~ εCR V / τ [ erg/s ].
17eVĖSNR?!E>1017eVĖ®;?!
E<10¹è­@, e-, p, He, …
=T,PQ,HñB
Ç,2012
Ø 宇宙線の滞在時間 τ が必要である.

=T,PQ,HñB Ç, 2012ÀýĠ100Uþ¯Ď宇宙線の起源について①

Ø CRs: 様々な核種で構成される
p 11B は星の元素合成でほとんど⽣成されないことが知られている．
11B production：
12C + p
ISM

→ 11B

※10B is also produced.
11B と 12C は銀河系内での宇宙線の振る舞いのトレーサーとなる．
Gabici+ 2019

Boron to Carbon ratio
0.4

Oliva+ 2017 (ICRC 2017)

0.3

B/C

0.2

0.1

0.05
0.04
0.03

C2/HEAO3
Webber et al.
CRN/Spacelab2
AMS01
ATIC02
CREAM-I
TRACERPAMELA
AMS02

0.02
1

10

102
EK [GeV/n]103

Boron to Carbon ratio
0.4

Oliva+ 2017 (ICRC 2017)

0.3

B/C

0.2

0.1

0.05
0.04
0.03
0.02

Solar abundance (Asplund+ 09)
10B : 11B = 2 : 8
-9.3
B
/
H
=
10
C2/HEAO3
Webber et al.
-3.6
C
/
H
=
10
CRN/Spacelab2
AMS01
B / C = 10-5.7
ATIC02(in number)
CREAM-I
TRACER
ほとんどのホウ素は
PAMELA
AMS02
12C + p11B→
ISM
で⽣成されたとして良い.
2
1

10

10
EK [GeV/n]103

宇宙線の起源について①
Ø 滞在時間の推定
11B production： 12C + p
11
ISM → B
⽣成率 (s-1)： q ~ n(12C) nISM σc
逃⾛率(s-1)： ~ n(11B) / τ
平衡状態：
n(11B) / τ = q

τ ~ {n(11B) / n(12C) } / (σ nISM c)＊Λ = μmp nISM τc is usually used in CR community.
→ n(11B) / n(12C) ~ σΛ / μmpτ ~ Λ / μmp nISM c~ 10 Myr (Λ / 10 g cm-2) (nISM / 1 cm-3)-1

宇宙線の起源について①
Ø CR は主に陽⼦で構成される

τ ~ Λ / μmp nISM c~ 10 Myr. (Λ / 10 g cm-2) (nISM / 1 cm-3)-1＊11B の滞在時間
Ø ISM は希薄なガス. → 宇宙線の集団的運動は電磁場によって決まる（粒⼦間衝突ではなく）．
Ø ジャイロ半径が宇宙線の運動の最⼩スケール.
→ rg ~ R/B, R = pc/Ze “rigidity”．
→ R が同じ粒⼦は同じ運動をする.

宇宙線の起源について①
Ø CR は主に陽⼦で構成される

τ ~ Λ / μmp nISM c~ 10 Myr. (Λ / 10 g cm-2) (nISM / 1 cm-3)-1
宇宙線が銀河を横切る時間: (10 kpc) / c ~ 0.01 Myr << τFinally escapingü CRは（おそらく）散乱により銀河内に閉じ込められている.
ü 拡散的に振る舞うことで，（おそらく）銀河内に充満しcf. Skilling 1971, 1975ている．

sourceεCRは“⼀様である”と仮定する．

宇宙線の起源について①
ü エネルギー密度
εCR~ 1 eV cm-3

ü 滞在時間τ ~ Λ / μmp nISM c~ 10 Myr (Λ / 10 g cm-2) (nISM / 1 cm-3)-1
ü 必要な注⼊率（Required Power）
PCR ~ εCR V / τ ~ εCR Mc / Λ~ 1041 erg/s ( εCR / 1 eV cm-3)( M / 109 Msolar )( Λ / 10 g cm-2 )-1＊ρ = μmpnISM ~ M /V is used.

宇宙線の起源について①
Required power: PCR ~ 1041 erg/sØ Supernova
ü Energy per event : ESN = 1051 erg
ü
Event rate: qSN ~ (Number of SNR) / (life time)~ 300/104 yr~ 1 / 30 yr→ Available power: ~ 1051 erg / 30 yr ~ 1042 erg/sØ 超新星爆発の運動エネルギーの〜１０％ほどが宇宙線の⽣成に消費されていれば良い（が，何故そうなるかは未解明）．
Other advantages:
1. Chemical composition of CRs (enhanced population of heavy elements)
2. Plausible CR acceleration mechanism was proposed (DSA)

宇宙線の起源について②“quasi-equilibrium” & “uniform εCR” are OK?
τ ~ 10 Myr. (Λ / 10 g cm-2) (nISM / 1 cm-3)-1
PCR ~ 1041 erg/s ( εCR / 1 eV cm-3) (M / 109 Msolar) (Λ / 10 g cm-2 )-1

ü 滞在時間のスケールは銀河の時間スケールに⽐べて⾮常に⼩さい．
e.g. 星形成率: ~ 1 Msolar / yr

ISMガスの質量 M ~ 109 Msolar は ~ 10 Myr
では変化しない．

宇宙線の起源について②“quasi-equilibrium” & “uniform εCR” are OK?
PCR ~ 1041 erg/s ( εCR / 1 eV cm-3) (M / 109 Msolar) (Λ / 10 g cm-2 )-1

Number of stars

ü~ 10 Myr の間に星形成率は~ 10 倍ほど変動する．
OB associationü時間的・空間的に宇宙線の注⼊率(by supernova)とεCR は⼀様とは限らない？
Palla & Stahler 2000

εCR ~ 1 eV cm-3 は銀河の代表値なのか？？？

宇宙線（Cosmic Rays）
CE²ĕīĽŐĮŔĶńįĻŐThe energy spectrum of CRs

CE²ĕīĽŐĮŔĶńį
ĽŐĮŔĕ
分⼦雲の電離度が局所的なOhira, Yamazaki & Terasawa (2012)εCRの⼿がかりになるだろう．

­@, e-,Ionizationp, He, …

ðīĽŐĮŔĕ
ýĠ100Uþ¯Ď¹è­@, e-,f p, He, …

forSuper Novae
e
g
n
g
i
Tr
tio
a
ÀýĠ100Uþ¯Ď
rm
17
o
f
0 eVĖSNR?!E>1017eVĖ®;?!

ðīĽŐĮŔĕ
17eVĖSNR?!E>1017eVĖ®;?!
E<10-, p, He, …
¹è­@,eMolecular=T,PQ,HñB

CloudÇ, 2012

=T,PQ,HñB Ç, 2012ÀýĠ100Uþ¯ĎFigure 6. Differential ionization rate of both proton and electron CRs (left-hand andISM spectra f0 assumed to be that from Voyager and AMS-02 fits. The black curves
ionization losses into the cloud.

分⼦雲の宇宙線電離率：観測 vs 理論
Phan+18

suc
ene
am
(
acc
vio
CR
(
com
tim
enh
pro
wh
(th
ene

分⼦雲の宇宙線電離率：観測vs 理論
CR ionization in diffuse molecular clouds

5171

Phan+18
CR protons

CR electrons

p Voyager 1 が太陽系“外”に進出したことで，外の宇宙
Figure 3. Data of the CR intensity for protons (left) and electrons (right) taken from Voyager 1 (Cummings et al. 2016) and AMS-02 (Aguilar et al. 2014
線量が初めて分かった．
2015) comparedwith the fitted curve used in this work.
p これが銀河の代表値と仮定すると，宇宙線電離率は≈ 10-17 s-1 is predictedとなり，観測を説明できない(Cummings+16, Phan+18).

Figure 6. Differential ionization rate of both proton and electron CRs (left-hand andISM spectra f0 assumed to be that from Voyager and AMS-02 fits. The black curves
ionization losses into the cloud.

分⼦雲の宇宙線電離率：観測 vs 理論
Phan+18

suc
ene
am
(
acc
vio
CR
(
com
tim
enh
pro
wh
(th
ene

Injection Power [erg s-1]

分⼦雲の宇宙線電離率：観測 vs 理論 Cosmic

1041 erg/s→ 10 % of SN power

Recchia+2019
E [eV]

Figure 1. Power needed in CR protons and electrons in orderto keep a carrot at a given energy in the whole Galactic disk,
able to predict (without taking into account the CR penetrationin MCs) an ionization rate of 4 ⇥ 10 16 s 1 , as compared with theelectron and proton energies in the range
より低いエネルギー
compare it with the total power (see e.g S
injected by sources in the observed CRの宇宙線(<MeV)
erg/s) and electron
spectrum で観(⇡ 1039 erg
an estimate of the total
CR proto-16 spower-1 in
測値≈4×10を説， suprathermin the whole Galactic disk theCummings et al. (2016) as
明しようとすると，
π 1M eV
4⇡J(E)E Vdisk
P
=
dE ⇡
C
2016
SNのほぼ全てのエネv p (E) ⌧loss, p (E)
1K eV
where J(E) is the CR proton flux of the
ルギーを低エネル(see Fig. 16 of Cummings et al. 2016).
Remarkably, the plot in Fig. 1 illust
ギー宇宙線の⽣成に
the short lifetime of low energy CRs in th
a CR carrot (or the suprathermal tail of
費やさないといけな2016) would require a power injection colarger (Recchia+2019).
than that already needed in orderい．
41

whole observed CR spectrum (⇡ 10 erg
is especially dramatic dramatic for electro
observed CR power for them is ⇡ 1039 erg
CRNote that 1041 erg/s roughly correspo
total power of galactic supernova explos
nova remnants are considered the major
CRs (see e.g Blasi 2013), our result imp
tence of a CR carrot would require eithe
large (in some cases even larger than 100* “uniform” εassumedξ energies=
fe (E)v
additional support to the idea that a CR carrot at
below
thee [1A(
Ẽ)
ẼV
energydetectedbyVoyager1).
Suchcomponent,
calledcarrot,disceI
smallest one
detected by Voyager 1 fails to provide a solution
to the (
.
P
(
Ẽ)
=
was first proposed by Meneguzzi, Audouze & Reeves (1971) toHere fp(e) (E) is the C
τlossofproblem(lightẼ) elements,of the ionizationin MCs.
a explain the abundancesand has rate
recentlybeen
v p(e) is the incident CR
reconsideredbyCummings
etal.
(2016,
whocalleditsuprathermalothe electron
σ ec isheight
Here Vdisc is the disc volume (radius Rd ∼ 15andkpc,
hdtail).
2 P OW E R R E Q U I R E M E N T
r Ø Coolingage secondary ionizatio
300
pc)
and
time
of
~ MeV
CRs and we analyse in
In this paper we focus on the
carrot scenariolino and/or& Gabici2015),
s detail the implications of the
Letpossible
us assume
the ofpresenceof a CR (electron
proton)presence
a CR population
potential.
4/3
1
component
at
a
given
energyẼ
!
3
MeV(energies
smaller thanIn
particular,
weestimate
the
power
at
energies
below
few
MeV.
(E)
≈
6
E
yr
for
E
in
1
keV–1
MeV
τ
e
loss,pkeV
Oncedetermined thedetected
uniformlythat has to be injected in those
low-energy
CRsbyinVoyager
order to 1),
keep
in the distributed in the whole
the power
e wholeτGalactic
2 ableGalactic
disc,
includingtheinteriorof MCs.tion
Forfunction,
simplicity,
we (
disc
a
populationtoaccountfortheobserved
(E)
≈
3
×
10
E
yr
for
E
in
1
keV–1
MeV
loss,ekeVin the whole Galactic diassumethatbytheassuming
distributionof such component is) ionization rate in MCs. We
do so
thatfunction
the carrot
A(Ẽ)ẼV
disc
component
is uniformly
distributedexpressions
both inside clouds
andthein powertheare theapproximate
forCR
protonandelectry Assumed
.
P
(
Ẽ)
=
Required
f
(E)
=
Aδ(E
−
Ẽ),
(1)
spectrumτloss (Ẽ)
rest of the ISM. The power estimated in that way represents aloss timeinthea Galactic
disc.
Such energy
losses ad very energy
conservative lower limit,
more realistic
scenario
wheresince
A is ainnormalization
constantthat has toHere
be determined.
Vdisc is the discmainlydue
toionization
lossesintheneutralphasesof thebyIS
low-energy
CRspresent
intheISM
penetrate
the
cloud
and
their
We do so by imposing that the H2 ionizationo 低エネルギー宇宙線はISMの荷電粒⼦とのクーロン散乱
300ratepc) producedandenergy lossesCRs
in MCshavebeintowithaccount.
As of the(electrons
or takenprotons)
the distributionfunctiongiven
by eandandCoulomb
lossesin totheionizedphasesISM(see
4/3
d transport
で冷える．
(E)
6 EkeV yr
shown by Phan et al. (2018),
the ionization
rate predictedin thiscaseξ ≈ 4 ×τloss,ps−1≈, detected
equation(1),equalsthe averagevalue,
10−16
Schlickeiser
ettheal.diffuse
2016).
a → would
be smaller than insimplescenario
presented
here. Hereinclouds(seee.g. Indriolo,
Fields & McCall
2009).
2
τ
(E)
≈
3
×
10
Eke
要求される注⼊エネルギーが増える.
loss,e5yr,seeequation5)
we showThethat, duetotherelatively
short
(!10
expressions of equation (5) are computed as
are the approximate ex
lifetime of sub MeV CRs in the ISM, in order to maintain a very
- low energy and hidden CR1component
energytime ofin thable to here
explainthe observed
We do1not consider
the effectof the CR carrot
on the loss
production= &, thea recent
τloss(p,e)
mainlydue &to Gabici
ionizatio(
lightelements.forFor
reviewof thistotopic see
Tatischeffrates,(E)it wouldbenecessarypotential
sources, ionization
and Coulomb losses in2018.
i ri(p,e)ftoi or larger than that needed
inject in the ISM a power comparable& to explain the whole observed CR spectrum. This result poses a
Schlickeiser et al. 2016
where ri and fi are the loss rate and filling factor, respectively, f

分⼦雲の宇宙線電離率：観測 vs 理論

分⼦雲の中に未知の源があるのだろうか？
(Yang+14; Recchia+19)

H2 line intensity in ~μm bands

分⼦雲の宇宙線電離率：
新しい測定法の提案 by VLT X-shooter
H2 + CR → H2* + CR

Bialy 2020

分⼦雲の宇宙線電離率：
新しい測定法の提案 by VLT X-shooter
宇宙線電離率が⾼いことは
H2 + CR→ H2* + CR
εCR > 1 eV/cc for ~GeV
CRsを⽰唆する？
H2 lines in ~μm bands

Ø 分⼦雲中で ~MeV CRsを作るのが必要なのか？
→具体的な天体・⽣成メカニズムは？
Ø SNが起源だと思って星形成率の増減で説明する？
→ 実際のところ，どれだけ⽣成されるのか？
Ø 分⼦雲の周りでの宇宙線の振る舞いは特殊？
→ 磁場が著しく乱れているといいのか？
Ø 分⼦雲の形成過程がヒント？
Bialy (arXiv 1910.12953)→ ガスを掃き集めるときに宇宙線も集める？

Supernova Remnant (SNR)
γ-ray: electron or proton?
X-ray: ~TeV CR electrons
Supernova ejecta

Ha: useful tracer of shockcondition & physics．
From Chandra archival image

SNR 0509-67.5 (Chandra & HST)
Blue: 1.5 – 7.0 keVGreen: 0.2 – 1.5 keV
Red: Ha

SNR shock は起源の第⼀候補である．

Supernova Remnant (SNR)
γ-ray: electron or proton? dFCR/dE ~ E-p e-E/Emax
pCR + pISM → π0 → 2γ
Inverse Compton
νFν d(log ν) ~ ν1-p
νFν d(log ν) ~ ν (3-p)/2
ν2Fν

p=2
From Chandra archival image

Uniform nISM is assumedlog νFν : photon energy flux
per frequency

Supernova Remnant (SNR)
γ-ray: electron or proton? dFCR/dE ~ E-p e-E/Emax

ガンマ線のスペクトルから…
1. 表⾯輝度から電⼦と陽⼦の総量が推定できる．
pCR + pISM → π0 → 2γ
InverseCompton
2. スペクトルの形から加速機構が制限される(スペク
1-p(3-p)/2
νFν d(logν)~ννFトル指数ν d(log ν)p~&ν最⾼エネルギー
E )
max

ν2Fν

p=2
From Chandra archival image

Uniform nISM is assumedlog νFν : photon energy flux
per frequency

Supernova Remnant (SNR)
γ-ray: electron or proton?

The Astrophysical Journal, 734:28 (9pp), 2011 June 10

E2dN/dE [ MeV cm-2 s -1 ]10-5

E2dN/dE [ MeV cm-2 s -1 ]10-5“電⼦” が⽀配的な形をしている.
→ 陽⼦ は加速してない？？？

Fermi LAT (24 months)
HESS (Aharonian et al. 2007)
Berezhko & Voelk 2006
Ellison et al. 2010 ( π 0dominated)
Zirakashvili & Aharonian 2010 (π 0 dominated)
Zirakashvili & Aharonian 2010 (IC/ π 0 mixed)

103

104

105
106
Energy [ MeV ](5) replacing the standard diffuse model by a model based on
the GALPROP code68 used in the Fermi-LAT analysis of theisotropic diffuse emission. The GALPROP model is describedin Abdo et al. (2010c). For (5), i.e., the GALPROP-based model,
we considered the various components of the diffuse emission
model separately for which we then individually fit the normalizations in our likelihood analysis. The components are gamma
rays produced by IC emission, gamma rays produced by interactions of CRs with atomic and ionized interstellar gas, andgamma rays produced in the interactions of CRs with moleculargas. The model component describing the gamma-ray intensityfrom interactions with molecular gas is further subdivided into
seven ranges of Galactocentric distance to accommodate localized variations of the CR and molecular gas density along theline of sight which are not accounted for in the model.
The same model of the isotropic component was used forall model variations (1)–(5). From model variations (1)–(5), we
obtain a systematic uncertainty of +0.08/−0.10 for the spectral index of RX J1713.7−3946 and a systematic uncertaintyof (+0.6/−0.7)×10−9 cm−2 s−1 for the flux above 1 GeV on
top of the statistical uncertainty. The systematic uncertainty ofthe derived flux and spectral index related to the uncertainty
in the LAT effective area was evaluated separately. The uncertainty of the LAT effective area—estimated from observationsof Vela (Abdo et al. 2009b) and the Earth albedo (Abdo et al.
2009a)—ranges from 10% at 500 MeV to 20% at !10 GeV.
The impact on the spectral parameters of RX J1713.7−3946 isa systematic uncertainty of ±0.05 for the spectral index and asystematic uncertainty of ±0.4 for the flux above 1 GeV. The
gray band in Figure 3 displays the superposition of all uncertainty bands obtained in our variations of the default model.
Figure 4 depicts the model variation (2) resulting in the softest
spectrum together with the fluxes in individual energy bands(black error bars) derived for model (2) using the same procedure as for the default model described above. The range ofsystematic uncertainty is particularly important to consider forcomparisons of the spectrum to pion-decay-dominated gammaray emission models which are generally expected to be softer
than IC-dominated gamma-ray emission models.

RX J1713.7-3946 の観測例

Abdo+201110-6

Abdo et al.

107

108

pCR + pISM → π0 → 2γ

From Chandra archival image ターゲットのガス密度が空間構造を持つ場合，電⼦の場合と似た形になり得ることが知られている(Inoue+12; Inoue 19).

Fermi LAT (24 months)

10-6

HESS (Aharonian et al. 2007)
Porter et al. 2006
Ellison et al. 2010 (IC dominated)
Zirakashvili & Aharonian 2010 (IC dominated)
3

10

4

10

105
106
Energy [ MeV ]

107

108

Supernova Remnant (SNR)
γ-ray: electron or proton?

The Astrophysical Journal, 734:28 (9pp), 2011 June 10

E2dN/dE [ MeV cm-2 s -1 ]10-5

Fermi LAT (24 months)
HESS (Aharonian et al. 2007)
Berezhko & Voelk 2006
Ellison et al. 2010 ( π 0dominated)
Zirakashvili & Aharonian 2010 (π 0 dominated)
Zirakashvili & Aharonian 2010 (IC/ π 0 mixed)

E2dN/dE [ MeV cm-2 s -1 ]

103

10-5

104

105
106
Energy [ MeV ]107(5) replacing the standard diffuse model by a model based on
the GALPROP code68 used in the Fermi-LAT analysis of theisotropic diffuse emission. The GALPROP model is describedin Abdo et al. (2010c). For (5), i.e., the GALPROP-based model,
we considered the various components of the diffuse emission
model separately for which we then individually fit the normalizations in our likelihood analysis. The components are gamma
rays produced by IC emission, gamma rays produced by interactions of CRs with atomic and ionized interstellar gas, andgamma rays produced in the interactions of CRs with moleculargas. The model component describing the gamma-ray intensityfrom interactions with molecular gas is further subdivided into
seven ranges of Galactocentric distance to accommodate localized variations of the CR and molecular gas density along theline of sight which are not accounted for in the model.
The same model of the isotropic component was used forall model variations (1)–(5). From model variations (1)–(5), we
obtain a systematic uncertainty of +0.08/−0.10 for the spectral index of RX J1713.7−3946 and a systematic uncertaintyof (+0.6/−0.7)×10−9 cm−2 s−1 for the flux above 1 GeV on
top of the statistical uncertainty. The systematic uncertainty ofthe derived flux and spectral index related to the uncertainty
in the LAT effective area was evaluated separately. The uncertainty of the LAT effective area—estimated from observationsof Vela (Abdo et al. 2009b) and the Earth albedo (Abdo et al.
2009a)—ranges from 10% at 500 MeV to 20% at !10 GeV.
The impact on the spectral parameters of RX J1713.7−3946 isa systematic uncertainty of ±0.05 for the spectral index and asystematic uncertainty of ±0.4 for the flux above 1 GeV. The
gray band in Figure 3 displays the superposition of all uncertainty bands obtained in our variations of the default model.
Figure 4 depicts the model variation (2) resulting in the softest
spectrum together with the fluxes in individual energy bands(black error bars) derived for model (2) using the same procedure as for the default model described above. The range ofsystematic uncertainty is particularly important to consider forcomparisons of the spectrum to pion-decay-dominated gammaray emission models which are generally expected to be softer
than IC-dominated gamma-ray emission models.

RX J1713.7-3946 というSNRの観測例”形”はガス密度の空間構造にも依存する．
→
加速機構の制限には不定性が存“電⼦”が⽀配的な形をしている.
在する．
→ 陽⼦ は加速してない？？？

Abdo+201110-6

Abdo et al.

108

pCR + pISM → π0 → 2γ

From Chandra archival image ターゲットのガス密度が空間構造を持つ場合，電⼦の場合と似た形になり得ることが知られている(Inoue+12; Inoue 19).

Fermi LAT (24 months)

10-6

HESS (Aharonian et al. 2007)
Porter et al. 2006
Ellison et al. 2010 (IC dominated)
Zirakashvili & Aharonian 2010 (IC dominated)
3

10

4

10

105
106
Energy [ MeV ]

107

108

Supernova Remnant (SNR)
γ-ray: electron or proton?
pCR + pISM → π0 → 2γ
Ø 表⾯輝度Lν,π0 ~ Fν ×4πd2 ~ 0.1 εCR V nISM σc”正確な” 距離 d と 密度 nISMが必要だが …
e.g.) RX J1713.7-3946
1. Caswell et al. 1975 : d ~ 10 kpc (HI absorption)
2. Fukui et al. 2003: d ~ 1 kpc (CO emission)
3. Wang et al. 2020: d ~ 5 kpc (red crump stars)(arXiv2005.08270)

Supernova Remnant (SNR)
γ-ray: electron or proton?

The Astrophysical Journal, 734:28 (9pp), 2011 June 10

E2dN/dE [ MeV cm-2 s -1 ]10-5

E2dN/dE [ MeV cm-2 s -1 ]

10

10-5

For ~ GeV CRs (Abdo+11)

Fermi LAT (24 months)
HESS (Aharonian et al. 2007)
Berezhko & Voelk 2006
Ellison et al. 2010 ( π 0dominated)
Zirakashvili & Aharonian 2010 (π 0 dominated)
Zirakashvili & Aharonian 2010 (IC/ π 0 mixed)
3

4

10

5

6

10
10
Energy [ MeV ]

7

10

8

10

From Chandra archival image

Fermi LAT (24 months)

10-6

HESS (Aharonian et al. 2007)
Porter et al. 2006
Ellison et al. 2010 (IC dominated)
Zirakashvili & Aharonian 2010 (IC dominated)
3

10

4

10

(5) replacing the standard diffuse model by a model based on
the GALPROP code68 used in the Fermi-LAT analysis of theisotropic diffuse emission. The GALPROP model is describedin Abdo et al. (2010c). For (5), i.e., the GALPROP-based model,
we considered the various components of the diffuse emission
model separately for which we then individually fit the normalizations in our likelihood analysis. The components are gamma
rays produced by IC emission, gamma rays produced by interactions of CRs with atomic and ionized interstellar gas, andgamma rays produced in the interactions of CRs with moleculargas. The model component describing the gamma-ray intensityfrom interactions with molecular gas is further subdivided into
seven ranges of Galactocentric distance to accommodate localized variations of the CR and molecular gas density along theline of sight which are not accounted for in the model.
The same model of the isotropic component was used forall model variations (1)–(5). From model variations (1)–(5), we
obtain a systematic uncertainty of +0.08/−0.10 for the spectral index of RX J1713.7−3946 and a systematic uncertaintyof (+0.6/−0.7)×10−9 cm−2 s−1 for the flux above 1 GeV on
top of the statistical uncertainty. The systematic uncertainty ofthe derived flux and spectral index related to the uncertainty
in the LAT effective area was evaluated separately. The uncertainty of the LAT effective area—estimated from observationsof Vela (Abdo et al. 2009b) and the Earth albedo (Abdo et al.
2009a)—ranges from 10% at 500 MeV to 20% at !10 GeV.
The impact on the spectral parameters of RX J1713.7−3946 isa systematic uncertainty of ±0.05 for the spectral index and asystematic uncertainty of ±0.4 for the flux above 1 GeV. The
gray band in Figure 3 displays the superposition of all uncertainty bands obtained in our variations of the default model.
Figure 4 depicts the model variation (2) resulting in the softest
spectrum together with the fluxes in individual energy bands(black error bars) derived for model (2) using the same procedure as for the default model described above. The range ofsystematic uncertainty is particularly important to consider forcomparisons of the spectrum to pion-decay-dominated gammaray emission models which are generally expected to be softer
than IC-dominated gamma-ray emission models.

RX J1713.7-3946 の観測例
(one of the brightest SNR?)

Abdo+201110-6

Abdo et al.

105106
Energy [ MeV ]

107

108

εCR V < 0.3 × 1051 erg(nISM/ 0.1 cm-3)-1 (d / 1 kpc)2※Proton origin is assumed.

Cosmic Rays
CE²ĕīĽŐĮŔĶńįĻŐThe energy spectrum of CRs

CE²ĕīĽŐĮŔĶńį
ĽŐĮŔĕOhira, Yamazaki & Terasawa (2012)­@, e-,Ionizationp, He, …

ðīĽŐĮŔĕ
ýĠ100Uþ¯Ď¹è­@, e-,f p, He, …

forSuper Novae
e
g
n
g
i
Tr
tio
a
ÀýĠ100Uþ¯Ď
rm
17
o
f
0 eVĖSNR?!E>1017eVĖ®;?!

ðīĽŐĮŔĕ
17eVĖSNR?!E>1017eVĖ®;?!
E<10-, p, He, …
¹è­@,eMolecular=T,PQ,HñB

CloudÇ, 2012

=T,PQ,HñB Ç, 2012ÀýĠ100Uþ¯Ď宇宙線（Cosmic
Rays）
宇宙物理学での宇宙線の例
CE²ĕīĽŐĮŔĶńįĻŐThe energy spectrum of CRs
1. ISMのエネルギー平衡 (太陽近傍)
Ohira, Yamazaki & Terasawa (2012)

CE²ĕīĽŐĮŔĶńį
CR ~ turbulence ~ B-field ~ thermal ~ 1 eV/cc
ĽŐĮŔĕ→ ガスの運動に影響を与える:­@, e-,Ionizationp, He, …
e.g. 銀河⾵の駆動源 (Breitschwerdt+ 1991)ðīĽŐĮŔĕ
2. 分⼦雲の電離源ýĠ100Uþ¯Ď¹è­@, e-,f p, He, …

forSuper Novaee
電離ガスは磁場と結合している．
g
n
g
i
Tr
tio
a
ÀýĠ100Uþ¯Ďrm
17o
磁場はプラズマ状態に依存して散逸する．
f0 eVĖSNR?!E>1017eVĖ®;?!

ðīĽŐĮŔĕ
17eVĖSNR?!E>1017eVĖ®;?!
→ 星形成中の磁束と⾓運動量の再分配(cf. Inutsuka 2012).
E<10-, p, He,
¹è­@,
e
…
Molecular=T,PQ,HñB
Ç,2012
宇宙線の起源は？

Cloud=T,PQ,HñB Ç, 2012ÀýĠ100Uþ¯Ď宇宙線（Cosmic
Rays）
宇宙物理学での宇宙線の例
CE²ĕīĽŐĮŔĶńįĻŐenergy CRsが担う)
spectrum of CRs
Ø εCR~1 eV cm-3 The( ~GeV

1. ISMのエネルギー平衡(太陽近傍)
Supernova
が最有⼒“候補“.
Ohira, Yamazaki & Terasawa (2012)
分⼦雲の電離度は星形成率で説明できるのか？
CR ~ turbulence ~ B-field ~ thermal ~ 1 eV/cc
ĽŐĮŔĕCR ~ turbulence ~ B-field ~ thermal ~ 1 eV/cc→ ガスの運動に影響を与える:-, p, He, …­@, e
は銀河(ISM)のどこでも成り⽴つのか？
Ionizatione.g.
銀河⾵の駆動源 (Breitschwerdt+ 1991)ðīĽŐĮŔĕ
分⼦雲の電離源
ýĠ100Uþ¯Ď
Ø2. SNRの先⾏研究¹è­@, e-,f p, He, …
forSuper Novaee
宇宙線の⽣成率と加速機構は未解明．
電離ガスは磁場と結合している．
g
n
g
i
Tr
tio
a
ÀýĠ100Uþ¯Ď
rm
17
17

CE²ĕīĽŐĮŔĶńį

磁場はプラズマ状態に依存して散逸する．
fo
0 SNRでの加速機構を解明するための提案：eVĖSNR?!E>10eVĖ®;?!
ðīĽŐĮŔĕ

17eVĖSNR?!E>1017eVĖ®;?!
星形成中の磁束と⾓運動量の再分配(cf. Inutsuka 2012).
E<10-, p, He,
¹è­@,
e
…
Hαの偏光観測が重要である. 2012
Molecular=T,PQ,HñBÇ,
宇宙線の起源は？
=T,PQ,HñB Ç, 2012
Cf. Shimoda etCloudal. 2018
Shimoda &ÀýĠ100Uþ¯ĎLaming 2019 a, b→

Diffusive Shock Acceleration; DSA
DSAの模式図

上流

(A)

下流

Assumption：① ”無衝突衝撃波”である．
②衝撃波を往復する⾼エネ
ルギー粒⼦が存在する．

磁気擾乱Alfven
wave

(B)
(C)

“power-law”分布を予⾔する．
dFCR/dE ~ E-p e-E/Emax

u0

衝撃波面

u2

p = 1 + 3u2 / (u0-u2)

Diffusive Shock Acceleration; DSA
Assumption：
DSAの模式図
p 背景の衝撃波構造 ( u0(r) & u2(r)) は宇宙線の反作⽤
上流
下流 ① ”無衝突衝撃波”である．
効果によって変調する(e.g. Drury & Volk 81; Drury &(A)②衝撃波を往復する⾼エネ
Falle 86)．
磁気擾乱
Alfven wave
ルギー粒⼦が存在する．
p

そのような反作⽤効果の証拠は⾒つ(B)
かっていない．
“power-law”分布を

(C)
予⾔する．
p 宇宙線がどれだけ加速されるかと，どのように振る
舞うか（≒拡散係数）を決めることが，宇宙線加速
dFCR/dE ~ E-p e-E/Emax
研究の最重要課題の⼀つである．

u0

衝撃波面

u2

p = 1 + 3u2 / (u0-u2)

宇宙線と衝撃波構造
no CR

Velocity

宇宙線なしの時の衝撃波構造(流体⼒学)

p: 背景プラズマの陽⼦

宇宙線が加速さ
れると？
up

down
shock (z=0)

z宇宙線と衝撃波構造
Velocityno CR

宇宙線がプラズマ
波動を励起し，上流は加熱される．

p

より多くの宇宙線が加速されると？

CR
heated/turbulent
up

precursor

down
shock (z=0)

z宇宙線と衝撃波構造
Velocity

no CR

宇宙線圧⼒が⼤きくなり，上流の流れを押し返す．→減速

p

p = 1 + 3u2 / (u0-u2)
with CRCR

up

precursor

down
shock (z=0)

宇宙線スペクトルは宇宙線⾃⾝が決める！

z宇宙線と衝撃波構造
Velocity

宇宙線圧⼒が⼤きくなり，上流の流れを押し返す．→減速
pp Cosmic Ray Modified Shock(CRMS)という．
no CR

p 上流の “加熱” precursorは光電離でも起p = 1 + 3u2 / (u0-u2)
きる (e.g. Ghavamian+00, Medina+14).
with CRCR

宇宙線スペクトルp 「速度変化」を確認してCRMSの証拠を
は宇宙線⾃⾝が決
める！
捉えたい．
up

precursor

down
shock (z=0)

zCosmic-Ray Modified Shock (CRMS)

010
No. 2, 2010

NONTHERMALRADIATION
OFOF
YOUNG
NONTHERMALRADIATIONYOUNGSNRs
SNRs

1 1

2
2
ρ/cρρ0 ,010P
ρ 0 Ṙg2 ,ρ10P
Ṙ2 ,c /10P
0 , 10P
0 Ṙ g /ρ 0 Ṙ
6 u, 103−1
10Te , keV
103 kms−1 ,km
10Ts e , ,keV//

πN(p)p(p)p4 cc //ρρ00Ṙ
Ṙff2
4 4πN
4

eam
tream
downstream downstream
reamupstrdownstream
downstream upstruepasm
RS

4RS
u
2

CD

CD
Te

FS

e

0.01

e

e

Pc

0.001

p

0.001
ρ

ρ
4

6

p/mc

p

Pg

Pg

8

r, pc
10 r, pc
12

pp

e

0.01

0

2

0.10.1

FS

Te

Pc

969969

0.01

1

100

p/mc
104

106

Figureof accelerated
0.013. Energy distributions
1
100 protons (thick
104 lines) and electrons106(multiplied to the factor of971
5000; thin lines) at the epoch t = 1620 yr. The
spectraat both theforward shock(solid lines)protons
and at thereverseshock
Figure3. Energy
distributionsof accelerated(thicklines)
and(dashedelectronslines) areto
shown.
(multiplied
the factor of 5000; thin lines) at the epoch t = 1620 yr. The

NONTHERMAL RADIATION OF YOUNG SNRs
Figure 2. Radial dependencies of the gas density (thick solid line), the gas
6
10 line), the gas12pressure
velocity (dottedline), the CR8 pressure(thick dashedSuzaku
◦
◦
◦
◦
◦
◦
◦
◦◦◦◦(thick(dasheddependenciesline), and the electron
temperatureline)line),
at t = the1620gasyr. The spectra at both the forward shock (solid lines) and at the reverse shock (dashed◦◦◦◦◦◦◦ solid
Radialof the gas
density(thin◦◦◦◦◦ the solid
calculations result
parameters ◦inpresent epoch: the forward lines) are shown.
2 in the following
−2 −1
log
F
(E)E
,
erg
cm
s
otted line), the CR pressure(thick dashed ◦◦◦line),
the gas pressure
EGRET
1
−10
shock
velocity 2760 km s−1 , its radius 10.5 pc,◦◦◦◦◦◦◦ the magnetic field strength
HESS
e),
and
the
electrontemperature
(thin
solid
line)
◦◦ at t = 1620 yr. The
downstream of the forward shock 127 µG. In the◦ same figure we show the ◦◦◦◦◦◦◦◦◦◦◦◦◦
s positions
result in ofthethefollowingparametersin the present
the πforward
4 πN (p)p4 c/ρ 0 Ṙf2
pforwardandreverse shocks,(FS and epoch:
RS, respectively)
and the ◦ ◦◦◦◦◦
0−1synch.
◦ ◦1
city
2760
km
s
,
its
radius10.5pc,
themagnetic
field
strength◦contact
discontinuity
(CD).
−11
◦◦◦
m of the forward shock 127 µG. In the same figure we show the
0.1
IC◦
4 πN (p)p4 c ρ 0 Ṙf2
p
f comparable
the forward andshocks,(FS
andRS,
respectively)
andtheto reverse
the ejecta
mass
at
t
>
100
yr
and
when
the
◦
TB

continuity(CD).
transition
to the Sedov phase begins. It is important that the−12Ø Zirakashvili & Aharonian 2010
先⾏研究のガンマ線スペクト/
ルの解釈では，宇宙線圧⼒がp
⾮常に⼤きい場合を考えてい
e
p
る． e

AT CA
observableratio Rb/Rf ∼ 0.5 is achieved at the Sedov phase π0 0.1◦the ejecta mass at t > 100 yr and when the
0.01
ble
to
when the FS speed is close to 1/3 of the characteristic ejecta−13n speedto theVejSedov
begins. ItX-ray
is important
the RX. Sincephasethe synchrotron
emission thatof SNR
leJ1713.7−3946
ratio Rb/Rfextends∼ 0.5wellis achieved
thetheSedov
beyond 1 at
keV,
presentphase
FS speed3−10.001
cannotbesignificantly
less
than
3
×
10
km
s
.
Therefore,e −14
FS speed is close to 1/3 of the characteristic ejecta log E, eV 0.01
−5
0
5
10
15
characteristic
velocity ofthe ejecta
cannotsmallerSince
thesynchrotron
X-ray
emission
ofbeSNR
RXthanj .the4−1
10 km
s . Suchvelocities
may
behadronic
attributed
to Ib/c−3946
well highbeyond
1 keV,within
the thepresent
FS speedodeling of extends
nonthermal radiationof RX J1713.7−3946
scenarioof gamma-ray production. The following basic parameters
f
b
−3
51
and
IIb
core
collapsesupernovae
withlowejectamasses,but 0.05, the electron-to-proton ratios at the e3−1= 1.2 kpc, nH = 0.09
cm than, ESN =32.7×× 10
Mej s= 1.5. MTherefore,⊙ , MA = MA = 23, ξ0 =0.001
eD significantly
less
10 erg,km
0.01
1
100
f
−4
b
−3

For RX J1713.7-3946

p/mc
104

e
106

Cosmic-Ray Modified Shock (CRMS)NONTHERMALRADIATION
OFOF
YOUNG
NONTHERMALRADIATIONYOUNGSNRs
SNRs

/

/

πN(p)p(p)p4 cc //ρρ00Ṙ
Ṙff2
4 4πN
4

eam
tream
downstream downstream
reamupstrdownstream
downstream upstruepasm
RS

4RS
u
2

CD

CD
Te

FS

ρ

ρ
4

6

pp

0.10.1
e

e

0.01
0.01
e 734:28 (9pp), 2011 June 10
The Astrophysical Journal,

e

Pc

0

2

FS

Te

Pc

9699691 1

22
ρ/cρρ0 ,010P
ρ 0 Ṙg2 ,ρ10P
Ṙ2 ,c /10P
0 , 10P
0 Ṙ g /ρ 0 Ṙ
6 u, 103−1
10Te , keV
103 kms−1 ,km
10Ts e , ,keVp0.001
Zirakashvili&
2010
p Aharonian
0.001
p/mc
Pg

Pg

r, pc
10 r, pc
128E2dN/dE [ MeV cm-2 s -1 ]010
No. 2, 2010

0.01

1

100

p/mc
104

106

Figureof accelerated
0.013. Energy distributions
1
100 protons (thick
104 lines) and electrons106(multiplied
5000; thin lines) at the epoch t = 1620 yr. The-5 to the factor of971
10
spectraat both theforward shock(solid lines)protons
and at thereverseshock
Figure3. Energy
distributionsof accelerated(thicklines)
and(dashedelectronslines) areto
shown.
(multiplied
the factor of 5000; thin lines) at the epoch t = 1620 yr. The

NONTHERMAL RADIATION OF YOUNG SNRs
Figure 2. Radial dependencies of the gas density (thick solid line), the gas
6
10 line), the gas12pressure
velocity (dottedline), the CR8 pressure(thick dashedSuzaku
◦
◦
◦
◦
◦
◦
◦
◦◦◦◦(thick(dasheddependenciesline), and the electron
temperatureline)line),
at t = the1620gasyr. The spectra at both the forward shock (solid lines) and at the reverse shock (dashed◦◦◦◦◦◦◦ solid
Radialof the gas
density(thin◦◦◦◦◦ the solid
calculations result
parameters ◦inpresent epoch: the forward lines) are shown.
2 in the following
−2 −1
log
F
(E)E
,
erg
cm
s
otted line), the CR pressure(thick dashed ◦◦◦line),
the gas pressure
EGRET
1
−10
shock
velocity 2760 km s−1 , its radius 10.5 pc,◦◦◦◦◦◦◦ the magnetic field strength
HESS
e),
and
the
electrontemperature
(thin
solid
line)
◦◦ at t = 1620 yr. The
downstream of the forward shock 127 µG. In the◦ same figure we show the ◦◦◦◦◦◦◦◦◦◦◦◦◦
s positions
result in ofthethefollowingparametersin the present
the πforward
4 πN (p)p4 c/ρ 0 Ṙf2
pforwardandreverse shocks,(FS and epoch:
RS, respectively)
and the ◦ ◦◦◦◦◦
0−1synch.
◦ ◦1
city
2760
km
s
,
its
radius10.5pc,
themagnetic
field
strength◦contact
discontinuity
(CD).
−11
◦◦◦
m of the forward shock 127 µG. In the same figure we show the
0.1
IC◦
4 πN (p)p4 c ρ 0 Ṙf2
p
f comparable
the forward andshocks,(FS
andRS,
respectively)
andtheto reverse
the ejecta
mass
at
t
>
100
yr
and
when
the
◦
TB
continuity(CD).
pFermi LAT (24 months)transition
to the Sedov phase begins. It is important that the
−12
-6
AT CA
HESS (Aharonian et al. 2007)
observableratio Rb/Rf ∼ 0.5 is achieved at the Sedov phase π0 0.1 10◦the ejecta mass at t > 100 yr and when the
0.01
ble
to
Berezhko & Voelk 2006
when the FS speed is close to 1/3 of the characteristic ejectae ( π 0dominated)
Ellison et al. 2010p−13
n speedto theVejSedov
phase
begins.
Itisimportantthatthe. Since the synchrotron X-ray emission of SNR RX
Zirakashvili & Aharonian 2010 (π 0 dominated)
leJ1713.7−3946
ratio Rb/Rfextends∼ 0.5wellis achieved
thetheSedov
beyond 1 at
keV,
presentphase
FS speede0.01
Zirakashvili & Aharonian 2010 (IC/ π 0 mixed)
be significantly
less than
3×
103 km s−1 . ejecta
Therefore,log E, eV 0.001
e cannotFS speedis close to 1/3of thecharacteristic−14

characteristic
velocity ofthe ejecta
cannotsmallerSince
thesynchrotron
X-ray
emission
ofbeSNR
RXthanj .the
4
−1
−5

0

5

10

Ø Zirakashvili & Aharonian 2010
Previous studies for the SED/
fitting supposed
the CR pressure
dominated case.

15

3

e

4

10
10
10 km
s . Suchvelocities
may
behadronic
attributed
to Ib/c−3946
well highbeyond
1 keV,within
the thepresent
FS speedodeling of extends
nonthermal radiationof RX J1713.7−3946
scenarioof gamma-ray production. The following basic parameters
e
f
b
−3
51
and
IIb
core
collapsesupernovae
withlowejectamasses,
but
3
−1
For
RX
J1713.7-3946= 1.2 kpc, nH = 0.09
cm than, ESN =32.7×× 10
Mej s= 1.5. MTherefore,
0.05, the electron-to-proton ratios at the⊙ , MA = MA = 23, ξ0 =0.001
eD significantly
less
10 erg,km
0.01
1
100
f
−4
b
−3

p/mc5
10
106
Energy4[ MeV ] 6
10
10

107

108

p

10

0

10

-1

10

-2

10

-3

10

-4

Cosmic-Ray Modified Shock (CRMS)• 太陽⾵の「その場観測」で確認済み．
太陽⾵
Vsh 〜 100 km/s
Vsw
ECR 〜 10 keV - MeV
速度変化 B〜 10 μG (MA 〜 5)
〜10 %Age 〜 day
Babs
若いSNR
Vsh >~ 1000 km/s
ECR 〜 1 GeV – 3 PeV ?
7
8
9
10
11
B
〜 1 - 100 μG ?
Terasawa 2006
(MA 〜 1 – 100 ?)~20 % of shock energy (flux)
Age 〜 100 – 1000 yr.
y of the upstream solar wind q was!130 km/s just before the arrival of the IP shock, whichconverted to non-thermal
particles
keV/cc1.5

energy density (EPIC 60-230 keV protons)

1.0
0.5

0.0
800

km/s

700
600
500

nT

400
10

5

0

4 event between 19 and 23 February, during which a moderately strong IP shock came at !09:03 UT on 21 February. Fromtained by the Goes particle monitor, energy density of energetic protons in the range of 60–230 keV !60–230, the solar wind
nitude of the magnetic field |B|. After the shock arrival VSW and |B| jumped to !900 km/s and J 20 nT, scaling out of theamounted to !1/4 of the total velocity jump across theCosmic-Ray Modified Shock (CRMS)

The Astrophysical Journal, 817:137 (5pp), 2016 February 1

Ohira

Ohi

ampliﬁed by the turbulence, which was suggested by Ohira(2013b). In this paper, we ﬁrst demonstrated via hybrid
simulations that even though shocks propagate partially ionizedplasmas, magnetic ﬁelds are ampliﬁed up to near theequipartition level by the downstream turbulence. Figure 1
shows that downstream magnetic ﬁelds are locally ampliﬁed
over 30B0 and the spatially averaged value is about 15B0 at
x  58, 400 c Xpp , that is, a signiﬁcant fraction of the
upstream kinetic energy is converted to the magnetic ﬁeld
energy. After the magnetic ﬁeld strength reaches a saturationlevel, it gradually decays as the turbulence decays. The decay
length scale is about 300 c Xpp , which is much smaller than theobservable length scale. However, once particles are accelerated by DSA, density ﬂuctuations with larger length scales
would be generated by the accelerated particles (Caprioli &The AstrophysicalJournal,817:137
1
Spitkovsky 2013).
If so, thedecay (5pp),
length 2016
scale Februaryof magnetic
Figure 2. Shock structures averaged over the y direction at t  2000 8cp1. T
ﬁelds would be larger. Although simulations in this paper are ared, blue, and black lines show the ﬂuid velocity normalized by the1 threetwo-dimensional
system,ithas
been
demonstratedby
theand normalized
the energy
by 10 spectru
times thevalue, uacceleration
x vd , the proton density
Figure 1. Two-dimensional shock structures at t  2000 8cp . From top upstream
to
upstream value, 0.1Sp Sp,0 , and the ionization fraction, fi , respectively.
dimensional magnetohydrodynamic simulation that the turbuin the previous simulation for vd bottom,lence
the panels
show proton density, Sp , magnetic
ﬁeld strength, ∣B∣, eachdoes amplify magnetic ﬁelds up
to near the equipartition
hand, the energy spectrum has a
Bx , By , Bactual
component
the magnetic
ﬁeld,Therefore,
ﬂuid velocity
of protons inz , andhigher
levelof(Inoueet al. 2010).
Alfvén Machsimulation for vd  1000 km s1
c Xpp ) are
he downstream
rest frame,
leftamplify
regions magnetic(x  57600
number shocks
can beﬁelds
ux , expectedu y . The tomorestrongly.
upstreamregions.
particle can experience the second
The amplitude of the magnetic ﬁeld ﬂuctuations is
The maximum velocity of leaking
EB B0 x 3 in the upstream region. The coherent length scale
3000 km s1 in the downstream resof magnetic ﬁeld structures is of the order of the gyroradius ofwith themagnetic
ﬁeld
strengthin theions
upstreamregion and the
upstream
pickupions.
The pickup
are preferentiallythe upstream region, they are picke
acceleratedbychargeexchange
andpickupprocesses(Ohirawave vector is highly oblique to the magnetic ﬁeld direction,and their velocity becomes about 4
2013b). Therefore, the observed magnetic turbulence isthat is,importantthe fastformagnetosonicmodes
are excited(Ohirarest frame. The picked up particles
injection into DSA.
Simulationsin thispaper 2014).
1, particle
performed infor
the vtwo-dimensional
system,km
so sthat
are heated to v x 8000 km s1, t
 30 v  1000
the correlaIn thisaresimulation

部分電離プラズマ中を伝播するSNR衝撃波のhybrid simulation(Ohira 16)．
Age 〜 1 dayCosmic-Ray Modified Shock (CRMS)

The Astrophysical Journal, 817:137 (5pp), 2016 February 1

Ohira

Ohi

ampliﬁed by the turbulence, which was suggested by Ohira(2013b). In this paper, we ﬁrst demonstrated via hybrid
simulations that even though shocks propagate partially ionizedplasmas, magnetic ﬁelds are ampliﬁed up to near theequipartition level by the downstream turbulence. Figure 1
shows that downstream magnetic ﬁelds are locally ampliﬁed
over 30B0 and the spatially averaged value is about 15B0 at
x  58, 400 c Xpp , that is, a signiﬁcant fraction of the
upstream kinetic energy is converted to the magnetic ﬁeld
energy. After the magnetic ﬁeld strength reaches a saturationlevel, it gradually decays as the turbulence decays. The decay
length scale is about 300 c Xpp , which is much smaller than theobservable length scale. However, once particles are accelerated by DSA, density ﬂuctuations with larger length scales
would be generated by the accelerated particles (Caprioli &The AstrophysicalJournal,817:137
1
Spitkovsky 2013).
If so, thedecay (5pp),
length 2016
scale Februaryof magnetic
Figure 2. Shock structures averaged over the y direction at t  2000 8cp1. T
ﬁelds would be larger. Although simulations in this paper are ared, blue, and black lines show the ﬂuid velocity normalized by the1 threetwo-dimensional
system,ithas
been
demonstratedby
theand normalized
the energy
by 10 spectru
times thevalue, uacceleration
x vd , the proton density
Figure 1. Two-dimensional shock structures at t  2000 8cp . From top upstream
to
upstream value, 0.1Sp Sp,0 , and the ionization fraction, fi , respectively.
dimensional magnetohydrodynamic simulation that the turbuin the previous simulation for vd bottom,lence
the panels
show proton density, Sp , magnetic
ﬁeld strength, ∣B∣, eachdoes amplify magnetic ﬁelds up
to near the equipartition
hand, the energy spectrum has a
Bx , By , Bactual
component
the magnetic
ﬁeld,Therefore,
ﬂuid velocity
of protons inz , andhigher
levelof(Inoueet al. 2010).
Alfvén Machsimulation for vd  1000 km s1
c Xpp ) are
he downstream
rest frame,
leftamplify
regions magnetic(x  57600
number shocks
can beﬁelds
ux , expectedu y . The tomorestrongly.
upstreamregions.
particle can experience the second
The amplitude of the magnetic ﬁeld ﬂuctuations is
The maximum velocity of leaking
EB B0 x 3 in the upstream region. The coherent length scale
3000 km s1 in the downstream resof magnetic ﬁeld structures is of the order of the gyroradius ofwith themagnetic
ﬁeld
strengthin theions
upstreamregion and the
upstream
pickupions.
The pickup
are preferentiallythe upstream region, they are picke
acceleratedbychargeexchange
andpickupprocesses(Ohirawave vector is highly oblique to the magnetic ﬁeld direction,and their velocity becomes about 4
2013b). Therefore, the observed magnetic turbulence isthat is,importantthe fastformagnetosonicmodes
are excited(Ohirarest frame. The picked up particles
injection into DSA.
Simulationsin thispaper 2014).
1, particle
performed infor
the vtwo-dimensional
system,km
so sthat
are heated to v x 8000 km s1, t
 30 v  1000
the correlaIn thisaresimulation

p ⽰唆される「速度変化」は〜１０％．
最低限，これを識別する必要がある．

部分電離プラズマ中をSNR(〜数kpc)にあるプラズマの〜１０％の伝播するSNR衝撃波
速度変化を判別せよ！
のhybrid simulation(Ohira 16)．
Age 〜 1 day

Challenging!

Supernova Remnant (SNR)
γ-ray: electron or proton?
X-ray: ~TeV CR electrons
Supernova ejecta

Ha: useful tracer of shockcondition & physics．
From Chandra archival image

SNR 0509-67.5 (Chandra & HST)
Haは衝撃波周辺のプラズマ状態を反映しBlue: 1.5 – 7.0 keV
SNR 91
shockて光る
forは起源の第⼀候補
review).
Green: 0.2 –(e.g.
1.5 keVRaymond
Red: Haである．

LEE ET AL.
Hα emission fromupstream

Hα Surface Brightness

Tycho’s SNR (Lee+2010)precursor

postshock

光電離？
Shock Direction

−8

−6

−4

−2

0

2

4

6

Distance from the peak [′′ ]le across the shock normal extracted from cut 04. The bright narrow peak is believed to correspond toLEE ET AL.
Hα emission fromupstream

Hα Surface Brightness

Tycho’s SNR (Lee+2010)precursor

postshock

光電離？
Shock Direction

The Astrophysical Journal Letters, 715:L146–L149, 2010 June 1

doi:10.1088/2041-8205/715/2/L146⃝
C 2010. The American Astronomical Society. All rights reserved. Printed in the U.S.A.

RESOLVED SHOCK STRUCTURE OF THE BALMER-DOMINATED FILAMENTS IN TYCHO’S SUPERNOVAREMNANT: COSMIC-RAY PRECURSOR?
Jae-Joon Lee1 , John C. Raymond2 , Sangwook Park1 , William P. Blair3 , Parviz Ghavamian4 , P. F. Winkler5 ,and Kelly Korreck2
1 Astronomy and Astrophysics Department, Pennsylvania State University, University Park, PA 16802, USA; lee@astro.psu.edu
2 Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA

3 Department of Physics and Astronomy, Johns Hopkins University, 3400 North Charles Street, Baltimore, MD 21218, USA
4 Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA
5 Department of Physics, Middlebury College, Middlebury, VT 05753, USA−8−6 Received 2010
−4February 24; accepted−2 2010 April 28;
0 published 20102May 1246Distancefrom the peak [′′ ]ABSTRACTWe report on the results from Hα imaging observations of the eastern limb of Tycho’s supernova remnant (SN1572)

le across the shock normal extracted from cut 04. The bright narrow peak is believed to correspond toKatsuda et al.
Hα emission fromupstream

Shock Direction

SNR Cygnus Loop
(Katsuda+2016)

Katsuda et al.
Hα emission fromupstream

Shock Direction
The Astrophysical Journal Letters, 819:L32 (6pp), 2016 March 10

SNR Cygnus Loop
(Katsuda+2016)
doi:10.3847/2041-8205/819/2/L32© 2016. The American Astronomical Society. All rights reserved.

SPATIALLY RESOLVED SPECTROSCOPY OF A BALMER-DOMINATED SHOCK IN THE CYGNUS LOOP:
AN EXTREMELY THIN COSMIC-RAY PRECURSOR?
Satoru Katsuda1, Keiichi Maeda2,3, Yutaka Ohira4, Yoichi Yatsu5, Koji Mori6, Wako Aoki7, Kumiko Morihana8,
John C. Raymond9, Parviz Ghavamian10, Jae-Joon Lee11, Jiro Shimoda4, and Ryo Yamazaki4
1

Institute of Space and Astronautical Science (ISAS), Japan Aerospace Exploration Agency (JAXA), 3-1-1 Yoshinodai,
Chuo, Sagamihara, Kanagawa 252-5210, Japan; katsuda@astro.isas.jaxa.jp
2
Department of Astronomy, Kyoto University, Kitashirakawa-Oiwake-cho, Sakyo-ku, Kyoto 606-8502, Japan
3
Kavli Institute for the Physics and Mathematics of the universe (WPI), The University of Tokyo, 5-1-5 Kashiwanoha, Kashiwa, Chiba 277-8583, Japan
4
Department of Physics and Mathematics, Aoyama Gakuin University, 5-10-1 Fuchinobe, Sagamihara 252-5258, Japan
5
Department of Physics, Tokyo Institute of Technology, 2-12-1 Ohokayama, Meguro, Tokyo 152-8551, Japan
6
Department of Applied Physics and Electronic Engineering, Faculty of Engineering, University of Miyazaki,
1-1 Gakuen Kibanadai-Nishi, Miyazaki 889-2192, Japan
7
National Astronomical Observatory of Japan, Mitaka, Tokyo 181-8588, Japan
8
Nishi-Harima Astronomical Observatory, Center for Astronomy, University of Hyogo, 407-2 Nichigaichi, Sayo-cho, Sayo, Hyogo 679-5313, Japan
9
Harvard-Smithsonian Center for Astrophysics, 60 Garden St., Cambridge, MA 02138, USA
10
Department of Physics, Astronomy and Geosciences, Towson University, Towson, MD 21252, USA
11
Korea Astronomy and Space Science Institute, Daejeon 305-348, Korea
Received 2016 January 28; accepted 2016 February 20; published 2016 March 7

ABSTRACTKatsuda et al.
Hα emission fromupstream

Shock Direction

SNR Cygnus Loop
(Katsuda+2016)

p Haは上流で光っている．
p 「速度変化」があると分かればCRMS．
p Hαの理論的な解釈は？

Hα emission from upstream
H

no CR

Velocityp

H はshock⾯を素通りする．
H + p/e → H* + p/eemits Lyβp/e はshock散逸を
受ける．

up

precursor

down fully ionizedshock (z=0)

zHα emission from upstream
H

no CR

Velocityp

Lyβ

H はshock⾯を素通りする．
H + p/e → H* + p/eemits Lyβp/e はshock散逸を
受ける．

up

precursor

down fully ionizedshock (z=0)

zHα emission from upstream
H

no CR

τLyβ ~ 1pup

Velocity

1s → 3p

Lyβ

H + p/e → H* + p/eemits Lyβp/e はshock散逸を
受ける．

Lyβ

precursor

H はshock⾯を素通りする．

down fully ionizedshock (z=0)

zHα emission from upstream
H

no CR

Pol. Hαup

τLyβ ~ 1p

Velocity

1s → 3p → 2s
Lyβ

Lyβ

上流Hα の起源：
1. Raman 散乱（Lyβ to Hα）．

Hαprecursor

down fully ionizedshock (z=0)

zHα emission from upstream
H

no CR

Pol. HατLyβ ~ 1p

Velocity

1s → 3p → 2s
Lyβ

Lyβ

上流Hα の起源：
1. Raman 散乱（Lyβ to Hα）．

Hα

y軸⽅向から観測したとすると，偏光⽅向は？
up

precursor

down fully ionizedshock (z=0)

zPolarization angle for Lyβ→Hα「Shock⾯に平⾏」な
直線偏光となる．
Lyβ→Hαdownstream
Scattered HaHα emission from upstream
H

no CR

Pol. Hαup

τLyβ ~ 1p

Velocity

1s → 3p → 2s
Lyβ

Lyβ

上流Hα の起源：
1. Raman 散乱（Lyβ to Hα）．

Hαprecursor

down fully ionizedshock (z=0)

zHα emission from upstream
H

no CR

VelocityPol. HατLyβ ~ 1pLyβ

H + e → H* + e

上流Hα の起源：
1. Raman 散乱（Lyβ to Hα）．
2. 「加熱」された電⼦による衝突励起．

unPol. Hαupprecursor

down fully ionizedshock (z=0)

zHα emission from upstream
H

no CR

Velocity

CR modified shockの場合．

pwith CRup

precursor

down fully ionizedshock (z=0)

zHα emission from upstream
H

no CR

Velocityp
H+p→p+H
with CRup

precursor「電荷交換反応」を介して，上流のHはpと同
じ速度構造を持ちうる．

down fully ionizedshock (z=0)

zHα emission from upstream
H

no CR

Velocityp

H はshock⾯を素通りする．
H + p/e → H* + p/eemits Lyβ

with CR

p/e はshock散逸を
受ける．

Lyβ
up

precursor

down fully ionizedshock (z=0)

zHα emission from upstream
H

no CR

Velocity

ドップラー効果により，この⽅向のLyβ は吸収されない．

pwith CR
Lyβup

precursor

down fully ionizedshock (z=0)

zHα emission from upstream
H

no CR

Velocityp
LyβPol. Hαwith CR

H + e → H* + eupprecursor

加熱された電⼦による衝突励起により，減速⽅向に垂直な⽅向に Lyβ が放射される．

Lyβ
down fully ionizedshock (z=0)

zPolarization angle for Lyβ→Hα「Shock⾯に垂直」
な直線偏光となる．
Lyβ→Hαdownstream
Scattered Ha
No scatteringin this direction.

Hα emission from upstream
H

no CR

Velocity

偏光⽅向と偏光度の定義:

Collisionalpexcitation in the「減速なし」→「shock⾯に平⾏」→偏光度＜０heating precursor
Lyβ
can also yield Lyβ「減速あり」→「shock⾯に垂直」→偏光度＞０perpendicularly toPol. Hαwith CR
the modification.

H + e → H* + eupprecursor

Lyβdown fully ionizedshock (z=0)

zHα emission from upstream

p Lineの輻射輸送とH原⼦の準位分布(rate-eq.)
Velocityno CR
H を偏光と同時に解く．→
極めて複雑！
偏光⽅向と偏光度の定義: This Lyβ photon isp
not scattered「減速なし」→「shock⾯に平⾏」→偏光度＜０p 簡単化して偏光度を評価する:because of the「減速あり」→「shock⾯に垂直」→偏光度＞０Doppler
shift.
1. 準位分布は偏光なしで解く (SJ& Laming
19a).
Stokes I はちゃんと計算できる．
with CR
2. Lyβ は完全に無偏光であるとする．
Lyβ

Pol. 3/2
Hα状態はLyβの吸収励起だけに由来
3. 上流の3p
Stokes Q を評価する.
precursor
up するとしてdown fully ionized
See, SJ & Laming 19b for detailsz
shock (z=0)

Cosmic ray modified shock polarimetry

Model set up for the shock

2727

Downloaded from https://academic

Shock下流は簡単のため，通常のランキン・ユゴニオ
Figure 3. Schematic illustration of the shock models considered in this関係で与え，T＝0.1
Tpofを仮定する.
paper.
The vertical axis showsethe
velocitieshydrogen atoms and protons
measured in the shock rest frame. The horizontal axis is the z-axis. The
shock (or sub-shock for the modified shock case) is located at z = 0. The

Cosmic ray modified shock polarimetry

Model set up for the shock

2727

We solve 3 cases:
1. No precursor （流体⼒学的な衝撃波）
Downloaded from https://academic

2. 電⼦加熱 あり, 減速なし(electron heating precursor)
3. 電⼦加熱 あり, 減速あり (Cosmic-Ray Modified Shock)

Shock下流は簡単のため，通常のランキン・ユゴニオ
Figure 3. Schematic illustration of the shock models considered in this関係で与え，T＝0.1
Tpofを仮定する.
paper.
The vertical axis showsethe
velocitieshydrogen atoms and protons
measured in the shock rest frame. The horizontal axis is the z-axis. The
shock (or sub-shock for the modified shock case) is located at z = 0. The

Line Transfer ModelLine transfer at SNR shock: 2s-state hydr

librium for the upstream plasma and fix theas T0 = 6000 K for simplicity.
Parameters:
Firstly, we describe the ionization structu
N be the number density of ’narrow’ hLet nH,j① 衝撃波速度Vshhave not experienced the charge-exchange② 上流の個数密度
ntot,0be the number densityof ’broad’ hydrogencharge-exchange reactions. Obviously, we ha③ 電離度（proton
fraction） χ0
N +nB . Fig. 1b is
nH,
schematic illustration ofH, jj④ 下流の陽⼦温度
Tdownof particles. We consider
that the partially ioni
the far upstream region (z < zout ) and presum⑤ 下流の電⼦温度
Te=βT
ionization equilibrium.
Hence,
we set the bo
down
N
N
B
nH,1s (z < 0) = nH,1s (zout ), nH,1s (z < 0) =np (zout ), where np is the number density of(z = 0), we assume the jump conditions as

Pure hydrogen plasma.
np (0) = 4np (zout ),励起準位は4fまで解かれている．
Vsh
u
=
,
2

(SJ & Laming 19a)k T

B p

=

Te

=

4
3 ′
µ mpVsh16
βTp

where Vsh is the shock velocity, kB is the Bo

2728

J. Shimoda and J. M. Laming

Results: Ionization Structure of H

Decelerated H

Results: Radiative vs. Collisional
上流の3p3/2状態はLyβの吸収励起だけに由来する. →だいたいokResults:
Polarizationof Ha
Cosmicray modified
shock polarimetry2729

The sign of degree indicates the polarization angle (Stokes Q).

Results: Polarization of Ha

Cosmic ray modified shock polarimetry

Polarization degree2729 2730

J. Shimoda and J. M. Laming

Surface brightness

Downloaded from https://academic.oup.com/mnras/article

偏光⽅向で「速度変化」の有無が分かる！
偏光度〜数％は観測できうる (Sparks+ 15).

Figure 6. The polarization degree of H α. The blue, red, and black lines Figure 7. The surface brightness profile of H α. The blue, red, and black
indicate the cases of no precursor (i), only an electron heating precursor lines indicate the cases of no precursor (i), only the electron heating(ii), and the electron heating precursor with decelerated protons (iii), precursor (ii), and the electron heating precursor with decelerated protons
respectively. The vertical thin lines at n tot, 0 z ≃ −27 × 1015 cm−2 and (iii), respectively. The vertical thin lines at ntot, 0 z ≃ −27 × 1015 cm−2 andn tot, 0 z ≃ 0 cm−2 indicate the locations of the precursor front and the shock ntot, 0 z ≃ 0 cm−2 indicate the locations of the precursor front and the shock
front, respectively. Positive values correspond to polarization along the z- front, respectively.
axis, while negative values correspond to polarization along the x-axis. Notethat these results includes all of atomic levels and processes unlike the result
is positive with a degree of ≃ 5 per cent. The degree of attenuation
shown in Fig. 2.

degree of H α, finally, is written as%# 1 (i) $ 3
9 2
N
dµ
I
−
µ
Q(s)
R
ν,µ
−1
16
16
ν
=

of Ly β photons radiated in the direction µ ≈ 1 depends on whether
the decelerated/non-decelerated atoms emit/absorb them, while thephotons radiated in the direction µ ≈ 0 are attenuated by the both
populations of atoms. Thus, the photons with µ ≈ 0 are efficiently
attenuated resulting in positive polarization.

The Astrophysical Journal Letters, 815:L9 (4pp), 2015 Decemb

Discovery of polarized Hα emission@ bright filament of SN 1006 (Sparks+ 15)
Figure 2. Upper, co-added spectra showing total intensity after basic reductions, and lower, same after subtraction of sky background showing strong Balmer
dominated shock spectrum. Lines visible are, right to left, Hα, Hβ, and Hγ.

Stokes I

Observed Area

Stokes Q

Stokes U

Ø Linear Polarization
Ø Polarization angle：
perpendicular to the shock
Figure 1. Location of 22 arcsec polarization slit segment on the Hubbl
TelescopeØ
(HST)
image of the SN 1006
North
is up, east is toDegree：2.0remnant.
± 0.4%
the SNR ﬁlament is in position angle ≈55°.

The Astrophysical Journal Letters, 815:L9 (4pp), 2015 Decemb

Discovery of polarized Hα emission@ bright filament of SN 1006 (Sparks+ 15)
Figure 2. Upper, co-added spectra showing total intensity after basic reductions, and lower, same after subtraction of sky background showing strong Balmer
dominated shock spectrum. Lines visible are, right to left, Hα, Hβ, and Hγ.

Stokes I

Observed Area

ü 観測と同程度の偏光度を推定できた．→理論
計算はよさそう．

ü Sparks+15の観測は空間分解されていないので，Stokes Q「速度変化」があるかはまだ分からない．
ü 上流・下流を空間分解した観測が必要．
Stokes U

Ø Linear Polarization
Ø Polarization angle：
perpendicular to the shock
Figure 1. Location of 22 arcsec polarization slit segment on the Hubbl
TelescopeØ
(HST)
image of the SN 1006
North
is up, east is toDegree：2.0remnant.
± 0.4%
the SNR ﬁlament is in position angle ≈55°.

Polarization of Ha vs others

Cosmic ray modified shock polarimetry

Polarization degree

Downloaded
Downloaded from https://academic.oup.com/mnras/article

Hβ/Hα2729 2730

J. Shimoda and J. M. Laming

Surface brightness

Figure 6. The polarization degree of H α. The blue, red, and black lines Figure 7. The surface brightness profile of H α. The blue, red, and black
indicate the cases of no precursor (i), only an electron heating precursor lines indicate the cases of no precursor (i), only the electron heating(ii), and the electron heating precursor with decelerated protons (iii), precursor (ii), and the electron heating precursor with decelerated protons
respectively. The vertical thin lines at n tot, 0 z ≃ −27 × 1015 cm−2 and (iii), respectively. The vertical thin lines at ntot, 0 z ≃ −27 × 1015 cm−2 andn tot, 0 z ≃ 0 cm−2 indicate the locations of the precursor front and the shock ntot, 0 z ≃ 0 cm−2 indicate the locations of the precursor front and the shock
front, respectively. Positive values correspond to polarization along the z- front, respectively.
axis, while negative values correspond to polarization along the x-axis. Notethat these results includes all of atomic levels and processes unlike the result
is positive with a degree of ≃ 5 per cent. The degree of attenuation
shown in Fig. 2.

degree of H α, finally, is written as%# 1 (i) $ 3
9 2
N
dµ
I
−
µ
Q(s)
R
ν,µ
−1
16
16
ν
=

「速度変化」を捉えているのは偏光観測だけ．

of Ly β photons radiated in the direction µ ≈ 1 depends on whether
the decelerated/non-decelerated atoms emit/absorb them, while thephotons radiated in the direction µ ≈ 0 are attenuated by the both
populations of atoms. Thus, the photons with µ ≈ 0 are efficiently
attenuated resulting in positive polarization.

decelerated
atoms emerging from the charge-exchange! 3(np /ntot,0 )Vshhydrogen/CCX , we obtain
ldiff
#
$2
reaction with25 thedecelerated
protons,
Vsh /CCX ("Vsh ), where "Vsh
V
sh
−2 −1
cm s
D(E) ! 1.05 × 10 −1
4000
km s−1
≃
200
km
s
is
the
relative
velocity of collision particles. In
Cosmic ray modified shock polarimetry2729#degree $−1
Polarization% n we&−1show≃ 3(np /ntot, 0 )Vsh /CCX . If
thispaper,
caseof zpre (35)
"Vthe
tot,0
sh
moda and× J. M. Laming,
1 cm−3200 km s−1

Outcome

Downloaded from https://academic.oup.com/mnras/article-abstrac

the modification comes from CRs accelerated at the shock, thelength-scale of modification can be given by their d#$ldiff ∼ D(E)/Vsh , where D and E are the CR diffu
"V
CCX
sh
−8
3 −1
≈ 4.59 × 10 cm s.
and energy, respectively. Thus, from the observat
np
200 km s−1
MNRAS 489, 2723–2731 (2019) ldiff ! 3(np /ntot,0 )Vsh /CCX , we obtain
The diffusion coefficient for the interstellar medium, DISM (1 GeV)#
$2
V
∼ 1028 cm2 s−1 , satisfies this condition. When magnetic distursh
D(E) ! 1.05 × 1025 cm−2 s−1bances are excited around an SNR shock, we cannot currently derive4000 km s−1
the relevant diffusion coefficient. If we suppose the simplest case,#$−1
D(E) ≃ ξ B rg (E)c/3, where rg is the gyro radius of CR and ξ B % ntot,0 &−1"Vsh×,> 1 6.isThethepolarization
gyro factor
byand
theblack
energyFiguredegree characterizedof H α. The blue, red,lines spectrum of
1 cm−3200 km s−1
where we approximate

Downloaded from https://academic.oup.com/mnras/article＊D(E)がボーム極限の場合に対応するエネルギーü 偏光観測から減速領域の⻑さが分かる．

indicatethe casesof no precursorwe(i), obtain
only an electron
heating
precursor
magnetic
disturbances,a lowerbound
for the CR energy(ii), and the electron heating precursor with decelerated protons (iii), where we approximate
E responsible for the velocity modification ofrespectively. The vertical thin lines at n tot, 0 z ≃ −27 × 1015 cm−2 and#$$ precursor#
n tot, 0 z ≃ 0 cm−2 indicate the#locations of thefront and the$shock
2
"V
C
CX
sh
B
Vsh along the z1
front, respectively. Positive
values correspondto polarization≈ 4.59 × 10−8 cm3 s−1
.
×
E
!
33
TeV
−1
axis, while negative values
correspondtopolarization
along
the
x-axis.
Note
−1
n
200
km
s
p
ξB
100 µG
4000 km sthat these results includes all of atomic levels and processes unlike the resulttion
structure
blue, red,shown
in Fig.%2. of the hydrogen# atoms. The$−1
The diffusion coefficient for the interstellar medium
&
−1
"V
n
tot,0
sh
ate the casesof no precursor (i), only the electron×.
cm2 s−1 , satisfies this condition. When m∼ 1028(36)
−3
−1
1 cm
200 km s

ü 「速度変化」を担う宇宙線の拡散係数がわかる．
→加速モデルの制限となる．

degreeH α, finally,is written
asandtheof electron
heating
precursor
with decelerated
bances are excited around an SNR shock, we cannot$ 3 the9number%# 1lines
ely.Note
The
solid
are
densityofhydrogen(i)2
thatat the diffusion coefficient. If we suppose theNdµ E is based on the coefficientIν,µ estimate− 16 µ ofQ(s)R −1thisthe relevant
16ν=
perienced a charge-exchange reaction in the precursor

1 ξB
Emax
B ! 1 mG
Z 41015.5 eV
!
"−1 !
"−1
RSNR
Vsh
Winkler et al.
×
10 pc
3000 km s−1

SN 1006

nal, 781:65 (18pp), 2014 February 1

Br

igh
tX
-ra
yl
ob
e

ü density ~ 0.3 cm-3
ü Vsh ~ 3000 km/s @$%# 1
32.2 9kpc
Q = CR
ILyβ,µ− µ2 dµ−1

55:00.0

50:00.0

45:00.0-41:40:00.0

η ~ 0.1 ? (this work)θ∼0-42:00:00.0

θ ∼ π/2

05:00.0

10:00.0

16 16(Raymond+ 07; Katsuda+12)µ = cos θLπ 0ηρVsh 2 3
∼ 0.1
R nσc
2
!
"3 !
"2 &
'2
R
V
n
sh
∼ 1033 η erg/s ×
3 pc
0.01c
0.3 cm−3

1 ξB
Emax
B ! 1 mG
Z 41015.5 eV
!
"−1 !
"−1
RSNR
Vsh
×
Winkler et al.
10 pc
3000 km s−1

SN 1006

al Journal, 781:65 (18pp), 2014 February 1-41:40:00.0

η ~ 0.1 ? (this work)

50:00.0

45:00.0

ü density ~ 0.3 cm-3
ü Vsh ~ 3000 km/s @$%# 1
32.2 9kpc
Q = CR
ILyβ,µ− µ2 dµ

55:00.0−1θ∼0-42:00:00.0

θ ∼ π/2

05:00.0

10:00.0

16 16(Raymond+ 07; Katsuda+12)µ = cos θLπ 0ηρVsh 2 3
∼ 0.1
R nσc
2
!
"3 !
"2 &
'2
R
V
n
sh
∼ 1033 η erg/s ×
3 pc
0.01c
0.3 cm−3

1 ξB
Emax
B ! 1 mG
Z 41015.5 eV
!
"−1 !
"−1
RSNR
Vsh
×
Winkler et al.
10 pc
3000 km s−1

SN 1006

al Journal, 781:65 (18pp), 2014 February 1-41:40:00.0

η ~ 0.1 ? (this work)−1

Condon+17 (Fermi-LAT)

05:00.0

10:00.0

16 16(Raymond+ 07; Katsuda+12)µ = cos θθ∼0θ ∼ π/2-42:00:00.0

55:00.0

50:00.0

45:00.0

ü density ~ 0.3 cm-3
ü Vsh ~ 3000 km/s @$%# 1
32.2 9kpc
Q = CR
ILyβ,µ− µ2 dµLπ 0ηρVsh 2 3
∼ 0.1
R nσc
2
!
"3 !
"2 &
'2
R
V
n
sh
∼ 1033 η erg/s ×
3 pc
0.01c
0.3 cm−3

10:00.0

05:00.0-42:00:00.0

55:00.0

50:00.0

45:00.0-41:40:00.0

1 ξB
Emax
B ! 1 mG
Z 41015.5 eV
The Astrophysical Journal, 851:100 (7pp), 2017 December!
20
"−1 !
"−1
RSNR
Vsh
scatteringof high-energy electrons.
×al Journal, 781:65 (18pp), 2014 February 1
Winkler et al.
10 pc
3000
km s−1 dominated. Howev
tobe leptonic

SN 1006

η efﬁcient~ 0.1 ?hadron(thisaccelerationwork) inspectral asymmetry visible in S
evidence in this respect. -3

ü density ~ 0.3 cm
The~ Fermi
Collaboratio
ü Vsh
3000LATkm/s@$ ongoing support%# 1
from a number
3that
9kpc
2.2
2
Q = CR
ILyβ,µ− haveµsupported
dµ both the develo−1

Condon+17 (Fermi-LAT)

16the LAT
16 as well as scientiﬁc data
(Raymond+
07; Katsuda+12)
National Aeronautics
and Spaceµ=cosθDepartment of Energy in the Unit

à l’Energieθ Atomiqueand the Cent
∼0
Scientiﬁque/Institut National de
Physiqueθ ∼desπ/2Particules in Fran
Italiana and the Istituto Nazionale
the Ministry of Education, Cul
Technology (MEXT), High Ene
Figure 3. Leptonic modeling for the ﬁve TeV shell-type remnants: RCW86
Organization (KEK) and Japan Aer(Red), RXJ1713.7−3946 (Black), RXJ0852.0−4622 (Blue), HESSJ1731(JAXA) in Japan, and the KAW
2 and TeV
−347 (Green) and SN1006 (Purple). Spectral points inηρV
the GeVsh
Swedish Research Council and th
bands are shown for HESSJ1731−347
of
Lπ0and SN1006.
∼ 0.1For the references
R3 nσceach modeling, see Acero et al. (2015b), from which the original
2 version of thisBoard in Sweden.
!
"3Additional!
"support
ﬁgure is extracted.
2&'2anforn scienceRVshphase is gratefully acknowledged fr∼ 1033 η erg/s ×
3 pc Astroﬁsica
0.01cin Italy0.3
andcm
the−3
CentreSummaryØεCR ~1 eV cm-3 ( ~GeV CRsが担う)• Supernova が最有⼒“候補“.
• 分⼦雲の電離度は星形成率で説明できるのか？
• CR ~ turbulence ~ B-field ~ thermal ~ 1 eV/cc
は銀河(ISM)のどこでも成り⽴つのか？

Ø

SNRの先⾏研究• 宇宙線の⽣成率と加速機構は未解明．

• SNRでの加速機構を解明するための提案：

Hαの偏光観測が重要である.

Cf. Shimoda et al. 2018
Shimoda & Laming 2019 a, b（⾒てね）