Upload
others
View
5
Download
0
Embed Size (px)
Citation preview
Gamma rays from molecular clouds(part 2)
Stefano GabiciAPC, Paris
Overview
Brief introduction on cosmic rays
The supernova remnant hypothesis for the origin of cosmic rays
Why gamma ray astronomy?
Why molecular clouds?
Part 1: Passive Molecular Clouds -> how to use MCs as cosmic ray
barometers
Part 2: Illuminated Molecular Clouds -> how to use MCs to locate CR
sources and constrain CR propagation
|----- Class #2 -----|
(1) Molecular clouds as cosmic ray barometers
Gamma rays from the Milky Way
dnγ(Eγ) =qγ(Eγ ,r )
4πr2d3r =
qγ4π
dr dΩ
The gamma ray emission from the Milky Way is mainly due to p-p interactions
photon flux at Earth
gamma ray emissivity
Gamma rays from the Milky Way
dnγ(Eγ) =qγ(Eγ ,r )
4πr2d3r =
qγ4π
dr dΩ
The gamma ray emission from the Milky Way is mainly due to p-p interactions
photon flux at Earth
gamma ray emissivity
dnγ
dΩ=
rmax
0dr
qγ4π
∝ rmax
0dr NCR(ECR, r) ngas(r)
Gamma rays from the Milky Way
dnγ(Eγ) =qγ(Eγ ,r )
4πr2d3r =
qγ4π
dr dΩ
The gamma ray emission from the Milky Way is mainly due to p-p interactions
photon flux at Earth
gamma ray emissivity
dnγ
dΩ=
rmax
0dr
qγ4π
∝ rmax
0dr NCR(ECR, r) ngas(r)
measured by EGRET, FERMI...
If we know (we do) the gas distribution in the Milky Way we can use gamma ray observations to constrain spatial variations of the CR intensity
Molecular Clouds as CR barometers: GeVs
<- FERMI sky = CR+gas
Molecular Clouds as CR barometers: GeVs
<- FERMI sky = CR+gas
Casa
nova
...SG
... e
t al
, 201
0
gas column density ->
CO data from NANTEN telescope
Gamma rays from the Milky WayNANTEN: CO (J=1-0)
-> tracer of H2
LAB HI Survey (Karberla et al 05)
ASSUMPTION: spatial homogeneity of CRs
Distance (kpc)110 1 10
)1 k
pc1
s2(p
hoto
ns c
m
1010
910
810 sum
HIH2
Casa
nova
et
al, 2
010
@1 GeV<--
Gamma rays from the Milky WayNANTEN: CO (J=1-0)
-> tracer of H2
LAB HI Survey (Karberla et al 05)
ASSUMPTION: spatial homogeneity of CRs
Distance (kpc)110 1 10
)1 k
pc1
s2(p
hoto
ns c
m
1010
910
810 sum
HIH2
Casa
nova
et
al, 2
010
@1 GeV<--
Gamma rays from the Milky WayNANTEN: CO (J=1-0)
-> tracer of H2
LAB HI Survey (Karberla et al 05)
ASSUMPTION: spatial homogeneity of CRs
2% 84% 14%
Distance (kpc)110 1 10
)1 k
pc1
s2(p
hoto
ns c
m
1010
910
810 sum
HIH2
Casa
nova
et
al, 2
010
@1 GeV<--
Gamma rays from the Milky WayNANTEN: CO (J=1-0)
-> tracer of H2
LAB HI Survey (Karberla et al 05)
ASSUMPTION: spatial homogeneity of CRs
2% 84% 14%
the main contribution to the gamma ray emission comes
from ~ 1 kpc. If the CR intensity is different there, we expect a different flux.
-> “Tomography” with gamma rays!
CR are, on large scales, homogeneously distributed in
the MW (EGRET)
Molecular Clouds as CR barometers: GeVs
LSR velocity (m/s)160 140 120 100 80 60 40 20 0 20
310×
(K)
BT
0
20
40
60
80
100
H I (b)
Abd
o et
al,
2009
Goul
d Be
lt
loca
l arm
Pers
eus
oute
r ar
m
100o < l < 145o -15o < b < 30o
Recent results from the FERMI collaboration
<- density along line of sight
Molecular Clouds as CR barometers: GeVs
LSR velocity (m/s)160 140 120 100 80 60 40 20 0 20
310×
(K)
BT
0
20
40
60
80
100
H I (b)
Abd
o et
al,
2009
Goul
d Be
lt
loca
l arm
Pers
eus
oute
r ar
m
100o < l < 145o -15o < b < 30o
Recent results from the FERMI collaboration
<- density along line of sight
R (kpc)0 2 4 6 8 10 12 14 16 18 20
)1 sr1
(sH
Iq
0
5
10
15
20
252710×
radial profile ofgamma ray emissivity
->
Molecular clouds at GeV energies: FERMI
coun
ts/pi
xel
0
5
10
15
20
25
30
35
40
45
50
coun
ts/pi
xel
0
5
10
15
20
25
30
35
40
45
50
° = -30b
° = -20b
° = -10b
° = 0b
° =
190
l
° =
200
l
° =
210
l
° =
220
l
° =
230
l
Mono. R2
Orion A
Orion B
Orion molecular clouds
Oku
mur
a et
al,
2009
Molecular clouds at GeV energies: FERMI
coun
ts/pi
xel
0
5
10
15
20
25
30
35
40
45
50
coun
ts/pi
xel
0
5
10
15
20
25
30
35
40
45
50
° = -30b
° = -20b
° = -10b
° = 0b
° =
190
l
° =
200
l
° =
210
l
° =
220
l
° =
230
l
Mono. R2
Orion A
Orion B
Orion molecular clouds
Oku
mur
a et
al,
2009
d ~ 400 pc+
CR spectrum from GALPROP(8% less than local CRs)
Molecular clouds at GeV energies: FERMI
coun
ts/pi
xel
0
5
10
15
20
25
30
35
40
45
50
coun
ts/pi
xel
0
5
10
15
20
25
30
35
40
45
50
° = -30b
° = -20b
° = -10b
° = 0b
° =
190
l
° =
200
l
° =
210
l
° =
220
l
° =
230
l
Mono. R2
Orion A
Orion B
Orion molecular clouds
Oku
mur
a et
al,
2009
Energy (MeV)210 310 410
)-1
MeV
-1 s-2
cm2 (M
eVE
/dNd2
E
-610
-510
-410
Orion A Total
400M 3 10× 4.8) ± 7.5 ±(80.6 /ndf = 16.7/102
Orion A Total
400M 3 10× 2.6) ± 5.2 ±(39.5 /ndf = 9.3/102
Pionic
Electron Bremsstrahlung
d ~ 400 pc+
CR spectrum from GALPROP(8% less than local CRs)
Molecular clouds at GeV energies: FERMI
coun
ts/pi
xel
0
5
10
15
20
25
30
35
40
45
50
coun
ts/pi
xel
0
5
10
15
20
25
30
35
40
45
50
° = -30b
° = -20b
° = -10b
° = 0b
° =
190
l
° =
200
l
° =
210
l
° =
220
l
° =
230
l
Mono. R2
Orion A
Orion B
Orion molecular clouds
Oku
mur
a et
al,
2009
Energy (MeV)210 310 410
)-1
MeV
-1 s-2
cm2 (M
eVE
/dNd2
E
-610
-510
-410
Orion A Total
400M 3 10× 4.8) ± 7.5 ±(80.6 /ndf = 16.7/102
Orion A Total
400M 3 10× 2.6) ± 5.2 ±(39.5 /ndf = 9.3/102
Pionic
Electron Bremsstrahlung
d ~ 400 pc+
CR spectrum from GALPROP(8% less than local CRs)
MA ∼ 8.1× 104M⊙
MB ∼ 4.0× 104M⊙
(2) Illuminated molecular clouds (MC/SNR associations)
Evolution of SuperNova RemnantsSN -> Explosion in a cold, uniform medium
ESN = 1051 E51 ergExplosion energy:
Mass of ejecta: Mej ≈ 1÷ 10 M⊙
Msw -> mass swept up by the shock
Msw << Mejif free expansion phase
Evolution of SuperNova RemnantsSN -> Explosion in a cold, uniform medium
ESN = 1051 E51 ergExplosion energy:
Mass of ejecta: Mej ≈ 1÷ 10 M⊙
Msw -> mass swept up by the shock
Msw << Mejif free expansion phase
ESN =1
2Mej v2 v = 109 E1/2
51
Mej
M⊙
1/2
cm/s
constant velocity
Evolution of SuperNova Remnantsthe free-expansion phase ends when: Mej ≈ Msw
Evolution of SuperNova Remnantsthe free-expansion phase ends when: Mej ≈ Msw
uniform medium with density: gas ≈ 1.7× 10−24 g
4 π
3R3
s gas = Mej Rs ≈ 2
Mej
M⊙
13
pc
shock radius
Evolution of SuperNova Remnantsthe free-expansion phase ends when: Mej ≈ Msw
uniform medium with density: gas ≈ 1.7× 10−24 g
4 π
3R3
s gas = Mej Rs ≈ 2
Mej
M⊙
13
pc
shock radius
duration of the free expansion phase: t ≈ Rs
v≈ 200
Mej
M⊙
− 16
yr
Evolution of SuperNova Remnants
Msw >> Mej -> the shock slows down
we want to find a relation between Rs and t
Evolution of SuperNova Remnants
Msw >> Mej -> the shock slows down
we want to find a relation between Rs and t
kbT2 =316
mu21 1 keV
let’s start by considering the shock heating of the gas
τc ∝ T12 106 yr
cooling time much longer than the SNR age!
Evolution of SuperNova Remnants
Msw >> Mej -> the shock slows down
we want to find a relation between Rs and t
SNRs emit X-rays
the SNR in this phase conserves the total energy!
kbT2 =316
mu21 1 keV
let’s start by considering the shock heating of the gas
τc ∝ T12 106 yr
cooling time much longer than the SNR age!
the only relevant physical quantities are: ESN gasand
Evolution of SuperNova Remnants
Msw >> Mej -> the shock slows down
we want to find a relation between Rs and t
the only relevant physical quantities are: ESN gasand
Evolution of SuperNova Remnants
Msw >> Mej -> the shock slows down
we want to find a relation between Rs and t
ESN
gas
t2
R5s
we can built a non dimensional quantity ->
Rs ≈ESN
gas
15
t25 Sedov solution
Evolution of SuperNova Remnantsduration of the Sedov phase
tage
Evolution of SuperNova Remnantsduration of the Sedov phase
tage vs ∝ t−35
Evolution of SuperNova Remnantsduration of the Sedov phase
tage vs ∝ t−35 T2 ∝ v2s
Evolution of SuperNova Remnantsduration of the Sedov phase
tage vs ∝ t−35 T2 ∝ v2s τc ∝ T
122
Evolution of SuperNova Remnantsduration of the Sedov phase
tage vs ∝ t−35 T2 ∝ v2s τc ∝ T
122
tage ∼ τcwhen -> radiative losses become important
tage ∼ 5× 104 yrthis happens at
Rends ≈ 20 pc
vends ≈ 200 km/s
Diffusive Shock Acceleration
ShockUp-stream Down-stream
u1u2
-->
B
B
B
B
B
B
Shock rest frame
Krymskii 1977, Axford et al. 1977, Blandford & Ostriker 1978, Bell 1978
See class by T. Bell
Diffusive Shock Acceleration
ShockUp-stream Down-stream-->
B
B
B
B
B
B
Up-stream rest frame
u1
u1 − u2
a
b
Ea = Eb
Diffusive Shock Acceleration
ShockUp-stream Down-stream-->
B
B
B
B
B
B
Down-stream rest frame
u1 − u2
a
b
Ea = Eb
u2
Particle acceleration during the free expansion phase
us
Up-stream
CR
shock
ld = us t
ld =√D t
Particle acceleration during the free expansion phase
us
Up-stream
CR
shock
ld = us t
ld =√D t
t =D
u2s
∝ Eequating...
Bohm diffusion -> λ = RL
D ≈ λc → D ∝ E
Particle acceleration during the free expansion phase
us
Up-stream
CR
shock
ld = us t
ld =√D t
t =D
u2s
∝ Eequating...
Bohm diffusion -> λ = RL
D ≈ λc → D ∝ E
max energy increases w. time
age max
Particles escape from SNRs
ush
Rsh
THIS IS A SNR
We need to know a bit of shock acceleration theory...
Diffusion length: ldiff ∼D(E)ush
∝ E
Bshush
see e.g. Ptuskin & Zirakashvili, 2003
Particles escape from SNRs
ush
Rsh
THIS IS A SNR
We need to know a bit of shock acceleration theory...
Diffusion length:
Confinement condition:
ldiff ∼D(E)ush
∝ E
Bshush
D(E)ush(t)
< Rsh(t) → Emax ∼ Bsh ush(t) Rsh(t)
see e.g. Ptuskin & Zirakashvili, 2003
Particles escape from SNRs
ush
Rsh
THIS IS A SNR
We need to know a bit of shock acceleration theory...
Diffusion length:
Confinement condition:
ldiff ∼D(E)ush
∝ E
Bshush
D(E)ush(t)
< Rsh(t) → Emax ∼ Bsh ush(t) Rsh(t)
Sedov phase:
Rsh(t) ∝ t2/5
ush(t) ∝ t−3/5
see e.g. Ptuskin & Zirakashvili, 2003
Particles escape from SNRs
ush
Rsh
THIS IS A SNR
We need to know a bit of shock acceleration theory...
Diffusion length:
Confinement condition:
ldiff ∼D(E)ush
∝ E
Bshush
D(E)ush(t)
< Rsh(t) → Emax ∼ Bsh ush(t) Rsh(t)
Sedov phase:
Rsh(t) ∝ t2/5
ush(t) ∝ t−3/5
Emax ∝ Bsht−1/5
see e.g. Ptuskin & Zirakashvili, 2003
Particles escape from SNRs
ush
Rsh
THIS IS A SNR
We need to know a bit of shock acceleration theory...
Diffusion length:
Confinement condition:
ldiff ∼D(E)ush
∝ E
Bshush
D(E)ush(t)
< Rsh(t) → Emax ∼ Bsh ush(t) Rsh(t)
Sedov phase:
Rsh(t) ∝ t2/5
ush(t) ∝ t−3/5
Emax ∝ Bsht−1/5
Emax decreases with timeParticles with E > Emax escape the SNR
Bsh also depends on
time
see e.g. Ptuskin & Zirakashvili, 2003
Particle escape from SNRsPtuskin & Zirakashvili, 2003
ush [km/s]
t [yr]Emax [GeV]
Particle escape from SNRsPtuskin & Zirakashvili, 2003
ush [km/s]
t [yr]Emax [GeV]
200 yrs
~ PeV
Particle escape from SNRsPtuskin & Zirakashvili, 2003
ush [km/s]
t [yr]Emax [GeV]
~10 GeV
~105 yrs
Particle escape from SNRsPtuskin & Zirakashvili, 2003
ush [km/s]
t [yr]Emax [GeV]
~10 GeV
~105 yrs
particles are released gradually during the Sedov phase: the one with the
highest energy first, and the ones with lower and lower energies at later times
The W28 region in gammas, CO, & radio
W28
HESS - TeV emission
The W28 region in gammas, CO, & radio
W28
HESS - TeV emission
Aharonian et al, 2008
The W28 region in gammas, CO, & radio
W28
18h18h03m
-24o
-23o 30
´-2
3o
0
20
40
60
80
100
NANTEN 12CO(J=1-0) 0-10 km/s
l=+6.0o
l=+7.0o
b=-1.0o
b=0.0o
W 28 (Radio Boundary)
W 28A2
G6.1-0.6
6.225-0.569
GRO J1801-2320
HESS J1801-233
HESS J1800-240A B
C
RA J2000.0 (hrs)
Dec
J20
00.0
(deg
)
18h18h03m
-24o
-23o 30
´-2
3o
0
20
40
60
80
100
NANTEN 12CO(J=1-0) 10-20 km/s
l=+6.0o
l=+7.0o
b=-1.0o
b=0.0o
W 28 (Radio Boundary)
W 28A2
G6.1-0.6
6.225-0.569
GRO J1801-2320
HESS J1801-233
HESS J1800-240A B
C
RA J2000.0 (hrs)
Dec
J20
00.0
(deg
)
NANTEN data - CO line
HESS - TeV emission
Aharonian et al, 2008
The W28 region in gammas, CO, & radio
W28
18h18h03m
-24o
-23o 30
´-2
3o
0
20
40
60
80
100
NANTEN 12CO(J=1-0) 0-10 km/s
l=+6.0o
l=+7.0o
b=-1.0o
b=0.0o
W 28 (Radio Boundary)
W 28A2
G6.1-0.6
6.225-0.569
GRO J1801-2320
HESS J1801-233
HESS J1800-240A B
C
RA J2000.0 (hrs)
Dec
J20
00.0
(deg
)
18h18h03m
-24o
-23o 30
´-2
3o
0
20
40
60
80
100
NANTEN 12CO(J=1-0) 10-20 km/s
l=+6.0o
l=+7.0o
b=-1.0o
b=0.0o
W 28 (Radio Boundary)
W 28A2
G6.1-0.6
6.225-0.569
GRO J1801-2320
HESS J1801-233
HESS J1800-240A B
C
RA J2000.0 (hrs)
Dec
J20
00.0
(deg
)
NANTEN data - CO line
HESS - TeV emission
Radio (Dubner et al 2000)
Aharonian et al, 2008
The W28 region in gammas, CO, & radio
W28
18h18h03m
-24o
-23o 30
´-2
3o
0
20
40
60
80
100
NANTEN 12CO(J=1-0) 0-10 km/s
l=+6.0o
l=+7.0o
b=-1.0o
b=0.0o
W 28 (Radio Boundary)
W 28A2
G6.1-0.6
6.225-0.569
GRO J1801-2320
HESS J1801-233
HESS J1800-240A B
C
RA J2000.0 (hrs)
Dec
J20
00.0
(deg
)
18h18h03m
-24o
-23o 30
´-2
3o
0
20
40
60
80
100
NANTEN 12CO(J=1-0) 10-20 km/s
l=+6.0o
l=+7.0o
b=-1.0o
b=0.0o
W 28 (Radio Boundary)
W 28A2
G6.1-0.6
6.225-0.569
GRO J1801-2320
HESS J1801-233
HESS J1800-240A B
C
RA J2000.0 (hrs)
Dec
J20
00.0
(deg
)
NANTEN data - CO line
HESS - TeV emission
Radio (Dubner et al 2000)
Aharonian et al, 2008
FERMI and AGILE - GeV emission
AGILE -> Giuliani et al, 2010
–23:00:00
–24:00:00
18:04:00 18:00:00
Right Ascension (J2000)
Dec
linat
ion
(J20
00)
0.1 0.2 0.3 0.4 0.5 0.6 0.7
HESS J1800-240
Source N
Source S
A B C
(a)
[counts/pixel]FERMI -> Abdo et al, 2010
W28: a cartoon
SNR shell
radio emission
-> GeV electrons
-> ongoing (re)acceleration?
W28: a cartoon
SNR shell
radio emission
-> GeV electrons
-> ongoing (re)acceleration?
Molecular clouds
W28: a cartoon
SNR shell
radio emission
-> GeV electrons
-> ongoing (re)acceleration?
Molecular clouds
gamma rays + gas
-> CR protons?
γ
γ
W28: a cartoon
SNR shell
radio emission
-> GeV electrons
-> ongoing (re)acceleration?
Molecular clouds
gamma rays + gas
-> CR protons?
-> coming from the SNR?
γ
γ
CR
CR
W28: a cartoon
SNR shell
radio emission
-> GeV electrons
-> ongoing (re)acceleration?
Molecular clouds
gamma rays + gas
-> CR protons?
-> coming from the SNR?
-> projection effects?
γ
γ
~12 pc
~ 20 pcCR
CR
?
W28: a cartoon
SNR shell
radio emission
-> GeV electrons
-> ongoing (re)acceleration?
Molecular clouds
gamma rays + gas
-> CR protons?
-> coming from the SNR?
-> projection effects?
+ masers
γ
γ
~12 pc
~ 20 pcCR
CR
?
W28: a cartoon
SNR shell
radio emission
-> GeV electrons
-> ongoing (re)acceleration?
Molecular clouds
gamma rays + gas
-> CR protons?
-> coming from the SNR?
-> projection effects?
+ masers
-> most of the protons
already escaped (> GeV?)
γ
γ
~12 pc
~ 20 pcCR
CR
Escaping particles assumptions
-> point like explosion
-> (isotropic) diffusion coefficient independent on position
-> particles with energy E escape after a time: tout ∝ E−δ
Escaping particles assumptions
-> point like explosion
-> (isotropic) diffusion coefficient independent on position
-> particles with energy E escape after a time: tout ∝ E−δ
this is just a parametrization
Escaping particles assumptions
-> point like explosion
-> (isotropic) diffusion coefficient independent on position
-> particles with energy E escape after a time: tout ∝ E−δ
higher energy particles are released first
when the SNR is smaller
this is just a parametrization
Escaping particles assumptions
-> point like explosion
-> (isotropic) diffusion coefficient independent on position
-> particles with energy E escape after a time: tout ∝ E−δ
higher energy particles are released first
when the SNR is smaller
this is just a parametrization
ld ≈ (D t)1/2
t = tage − tout
diffusion length
? model dependent...
Escaping particles assumptions
-> point like explosion
-> (isotropic) diffusion coefficient independent on position
-> particles with energy E escape after a time: tout ∝ E−δ
higher energy particles are released first
when the SNR is smaller
this is just a parametrization
ld ≈ (D t)1/2
t = tage − tout
diffusion length
? model dependent...
≈ (D tage)1/2
high energy particles
Escaping particles assumptions
-> point like explosion
-> (isotropic) diffusion coefficient independent on position
-> particles with energy E escape after a time: tout ∝ E−δ
higher energy particles are released first
when the SNR is smaller
this is just a parametrization
ld ≈ (D t)1/2
t = tage − tout
diffusion length
? model dependent...
≈ (D tage)1/2
high energy particles
lower energies -> point-like approx not so good + model dependent (tout ~ tage)
Escaping particles: CR spectrumtage -> ld = 10, 30, 100 pc
ld ≈ (D tage)1/2
Escaping particles: CR spectrumtage -> ld = 10, 30, 100 pc
const ∝ ηCR ESN
l3d
ld ≈ (D tage)1/2
Escaping particles: CR spectrumtage -> ld = 10, 30, 100 pc
const ∝ ηCR ESN
l3d
∝ e−
Rld
2
ld ≈ (D tage)1/2
Escaping particles: CR spectrumtage -> ld = 10, 30, 100 pc
const ∝ ηCR ESN
l3d
∝ e−
Rld
2
γ
ld ≈ (D tage)1/2
Escaping particles: CR spectrumtage -> ld = 10, 30, 100 pc
const ∝ ηCR ESN
l3d
∝ e−
Rld
2
γ
ld ≈ (D tage)1/2
0.1 ηCR 1
tage ≈ 3× 104 ÷ 105yr
ESN 2÷ 4× 1050erg
estimate D?
MNcl ≈ 5× 104M⊙
if we want SNRs to be the sources of CRs...
North
South
Fγ
Mcl
North
≈
Fγ
Mcl
South
Escaping particles: W28
North
South
Fγ
Mcl
North
≈
Fγ
Mcl
South
ld both clouds within a distance ld
Escaping particles: W28
North
South
Fγ
Mcl
North
≈
Fγ
Mcl
South
ld both clouds within a distance ld
Fγ ∝ fCR Mcl
d2
Escaping particles: W28
North
South
Fγ
Mcl
North
≈
Fγ
Mcl
South
ld both clouds within a distance ld
Fγ ∝ fCR Mcl
d2∝ ηCR ESN
(D tage)3/2
Mcl
d2
Escaping particles: W28
North
South
Fγ
Mcl
North
≈
Fγ
Mcl
South
ld both clouds within a distance ld
Fγ ∝ fCR Mcl
d2∝ ηCR ESN
(D tage)3/2
Mcl
d2
Escaping particles: W28
D ∝ηCRESN Mcl
Fγd2
2/3
t−1ageand we get... ->
W28: numbers
Measured quantities:
apparent size: ~ 45’-50’, distance ~ 2 kpc -> Rs ~ 12 pc ratio OIII/Hβ -> vs = 80 km/s
Assumptions:
mass of ejecta -> Mej ~ 1.4 Msun
ambient density -> n ~ 5 cm-3
Derived quantities (Cioffi et al 1989 model):
explosion energy -> ESN ~ 0.4 x 1051 erg initial velocity -> v0 ~ 5500 km/s SNR age -> 4.4 x 104 yr (radiative phase)
W28: TeV emission
North
South
A B
~ 5 104 Msun
~ 6 104 Msun
~ 4 104 Msun
Gabici et al, 2010
W28: TeV emission
North
South
A B
~ 5 104 Msun
~ 6 104 Msun
~ 4 104 Msun
North
South A
South B
virtually the same curves for:
(η,χ) = (10%, 0.028); (30%, 0.06); (50%, 0.085)
acceleration efficiency
diffusion coefficient in units of the galactic one
Gabici et al, 2010
W28: TeV emission
North
South
A B
~ 5 104 Msun
~ 6 104 Msun
~ 4 104 Msun
North
South A
South B
virtually the same curves for:
(η,χ) = (10%, 0.028); (30%, 0.06); (50%, 0.085)
acceleration efficiency
diffusion coefficient in units of the galactic one
diffusion
coe
ffici
ent is
supp
ressed
?
Gabici et al, 2010
W28: TeV emission
North
South
A B
~ 5 104 Msun
~ 6 104 Msun
~ 4 104 Msun
North
South A
South B
virtually the same curves for:
(η,χ) = (10%, 0.028); (30%, 0.06); (50%, 0.085)
acceleration efficiency
diffusion coefficient in units of the galactic one
size of “CR bubble” @3 TeV
ld ≈ 70 pc → χ = 0.028
≈ 95 pc → χ = 0.06
≈ 115 pc → χ = 0.085diffusion
coe
ffici
ent is
supp
ressed
?
Gabici et al, 2010
W28: GeV emission
Energy (GeV)110 1 10 210 310
) 1 s2
dF/d
E (e
V cm
2 E
110
1
10
210(b) W28 South (Source S)
Energy (GeV)110 1 10 210 310
) 1 s2
dF/d
E (e
V cm
2 E
110
1
10
(a) HESS J1800 240A
Energy (GeV)110 1 10 210 310
) 1 s2
dF/d
E (e
V cm
2 E110
1
10
210(a) W28 North (Source N)
FERM
I ->
Abd
o et
al,
2010
if the cloud is within the “CR bubble” [d < ld(E)]
if cloud A and B are at the same distance -> same spectrum
AB
cloud A might not be physically associated with the
W28 system
W28: GeV emission
W28: GeV emission
~ TeV
W28: GeV emission
~ TeV
~ GeV
W28: GeV emission
~ TeV
~ GeV
W28: GeV emission
d = 12 pc d = 32 pc d = 65 pc
η = 30%
χ = 0.06
North South B South A
W28: GeV emission
d = 12 pc d = 32 pc d = 65 pc
η = 30%
χ = 0.06this peak is removed as background
North South B South A
W28: GeV emission
d = 12 pc d = 32 pc d = 65 pc
η = 30%
χ = 0.06this peak is removed as background
some problems here...+ some of the approximations are not very good at low energies
+ subtraction of the background?+ other contributions? (Bremsstrahlung)
North South B South A
Take-home message
Current observations of MCs in gamma rays can be used to probe the CR
intensity at different locations within the Galaxy (i.e. Gal centre ridge, orion cloud,
W28) -> indirect identification of CR sources
For MC located close to the (supposed) CR accelerator, estimates of the local CR
diffusion coefficient can be made.
Future observations (Fermi, HESS/MAGIC/VERITAS, but most of all CTA) will
increase the number of detected clouds and make the results presented here more
(or less?) solid