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Pulsar wind nebulaePulsar wind nebulae
Isabelle A. GrenierUniversité Paris 7 & CEA Saclay
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useful bibliographyuseful bibliography
Gaensler & Slane, astro-ph/0601081 (with many references to observations)Arons J. 2004, Adv. in Space Research 33, 466Rees & Gunn 1974, MNRAS 167, 1Kennel & Coroniti 1984, ApJ 283, 694 and ApJ 283, 710Chevalier, R. 1982, ApJ 258, 790Romanova, Chulsky & Lovelace 2005, ApJ 630, 1020Wilkin 1996, ApJ 459, L31Bucciantini et al. 2003, A&A 405, 617Blondin, Chevalier, Frierson 2001, ApJ 563, 806van der Swaluw et al. 2003, A&A 397, 913de Jager et al. 1996, ApJ 457, 253Aharonian et al. for HESS, 2005, A&A 448, L43 and A&A 432, L25 and A&A 435, L17 and A&A 442, L25
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synchrotron nebula in a supernova remnantsynchrotron nebula in a supernova remnant
SNR G292.0+1.8 + PSR J1124-5916P = 135 ms, age = 1.7 kyr
SNR emission synchrotron nebula
0.6-2 keV 2-7 keV
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plerionplerion
optical Crab nebulablue synchrotron rad.wind nebula fillsthe remnant (plerion)
pulsar
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bow shock pulsar wind nebulabow shock pulsar wind nebula
PSR 1957+20 « Black Widow »recycled ms pulsarsupersonic motion in the ISMbow shock confines the windnebulaD ~ 1.5 kpc
40 to 50 PWN knownin the Milky Way + LMC + SMC
from most energetic youngpulsars (bias)
4 1029 W ≤ Ėpsr ≤ 5 1031 W“round” PWNe for young pulsars
τpsr ≤ 20 kyr“bow shock” PWNe for older pulsars
CXO opt
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Ω Bαlight
cylinderRLC= c/Ω
closedfield lines
E⊥B
nullcharge surfaceΩ.B= 0
polar capγ+B→ e±
outergap
γ+X→ e±
slot gap
pulsar magnetospherepulsar magnetosphereunipolar inductor field from B rotation
extracts charges from the neutron starmgnsph. filled with force-free charge density
Goldreich & Julian ’69acceleration sites in openmgnsph. where
primary charges accelerated to ~10 TeVemit γ photon (curvature rad.)
γ initiate e± cascades
question: are ions alsoaccelerated?
//GJ E⇒ρ≠ρ
B)r(Errrr
∧∧Ω−=
BGJvv⋅Ωε−=ρ 02
2psr m kg for E 3810≈ΩΩ= I&& I
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pulsar windpulsar windretarded potentials => elmgn. wave outside the light cylinderwind = toroidal (wound up) field frozen to an e±- ion plasmaσ ratio = Poynting flux / particle energy flux (Rees & Gunn ’74, Kennel & Coroniti ’84)
σ >> 1 EM wave takes the pulsar energy away (near light cylinder)σ << 1 pulsar energy transferred to wind particles (near terminal shock)
light cylinder
wind γ ∼106
ambiant medium
shocked wind(syn. neb)
contact discontinuity
psw= pext
terminal shockterminal shockterminal shock
B⊥Ω
)mc(vncB
www
w2
111
021
γ
μ=σ
σ+=
σ+=⇒+=
/E
EE
EEEE psrEB
psrwEBwpsr 111
&&
&&&&&
1>>σ 1>>σ
1<<σ
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open questionsopen questions
how are particles accelerated in the magnetosphere?how many secondaries produced in the cascades?
how is the elmgn energy transferred to the particles in the wind?σ >> 1 at the light cylinder, σ = 0.001 - 0.1 at the terminal shock‘silent’ transfer since charges follow B ⇒ syn. radiation <<large equatorial to pole asymmetries (torii and jets)
how is this energy transformed to radiation in the shocked wind?physics of a relativistic shock
how does the wind nebula evolve?early expansion in ‘cold’ SN ejectalater expansion in ‘hot’ ejecta shocked by the SNR reverse shockbow shock confinement from pulsar supersonic motion in SNR or in ISM
applicationsblazar and γ-ray bursts jet formation and relativistic shocksSag A*
« The answer, my friend, is blowing in the wind » (Bob Dylan)
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PWN structurePWN structure
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MHD shock waveMHD shock waveconservations => jump conditions in the shock frame
strong perpendicular shock p2 >> p1: compression ratio X
X 4 if pB 0 (Rankin-Hugoniot)X 0 if large pB
{ }{ } 012
41222
1111
11123
1
=+γ−++γ+
+−γ+γ−−γ
ramBramth
BramthB
p)(pppX)pp)((pXX)(p
[ ] [ ]
cteB andcteuBuB :fluxmagnetic )(
cteBuBuBupu)uu( :energy (4)
cteBBuu andcte)BB(up:momentum (2)
cteu:mass (1)
////
//////int//
//////
==−
=−μ
−+++ρ
=μ
−ρ=−μ
+ρ+
=ρ
⊥⊥⊥
⊥⊥⊥⊥
⊥⊥⊥⊥
⊥
5
1
12
1
0
2221
0
22
0
2
1
2
1
2
2
1BB
uuX =
ρρ
==⊥
⊥
⊥
//
u2 = u1/X u1
pram2 = ρ1u12/X p1 ~ ρ1u1
2
pB2 = X2pB1 pth2
B2 = XB1 ρ2 = Xρ1
Ma2 < 1 Ma1 >> 1
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relativistic shock waverelativistic shock wave
(*) = relativistic wind rest frameupstream = 1, downstream = 2
Kennel & Coroniti ’84, Blandford & McKee ’76, Double ’04
jump conditions
compression ratio:perpendicular case
if σ << 1 Y 3 modest shock (lower for oblique B)if σ >> 1 Y 1+(1/2σ) hardly any shock and thermalization of incident ram pressure
cte)Bpe( :energycteB:field
cteBp)Bpe( :momentumcten :mass
***
****
=μ
++βγ=γβ
=μ
++μ
++βγ=γβ
0
22
0
2
0
222
2
21
mcnpe :energy internal
pp :pressure
nnn :density
BBB :field
*
ad
*
*
*
*
+−γ
=
=
γ==
γ==
⊕
⊕
),(fvv
nn
BBY 11
2
1
2
1
1
2
1
2 γσ==ββ
===
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spherical PWN σ << 1spherical PWN σ << 1
SNR
PWN
wind
RTSRext
Requipartition
ppwn≈ pSNR
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spherical PWN σ <<1spherical PWN σ <<1
density, velocity, B profiles
wind PWN
pth > pB Requip pB > pth
pth≈ cte pB≈ ctecRE)(p
TS
wramth 23
132
224π
+≈+
&
Rees & Gunn ‘74Kennel & Coroniti ’84
ptot = cte = pext v∞ ≠ 0 to carry B flux
12
0
2
2
2
2
0
0
0
34
−φ
−∞
φ
φ
−
γ−
∝∝⇒
σ≈→
≈μ
∝⇒=
=∧+
∝⇒
==
=⇒⇒
=γ=
rB,rn
cvv:velocityasymptotic
nkTB
for R
rB)Bvr(drd
)Bv(rot :ystationnarMHD
rv
)nvr(drd)vn(div :conserved mass
cten isobaric subsonic /,ctepn
r
ionequipartit
r
adad
rr
r
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σ >> 1 case: no PWNσ >> 1 case: no PWNdensity, velocity, B profiles
elmgn wave dominatesgoes onpth2 << pB2 ⇒ little e± dispersion and little synchrotronpB2/pram2= 2σ et pram2/pth2 = 8σ²
windpB >> pth
cr
EB psr2
0
2
42 π→
μ∞
&
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spherical PWN σ << 1spherical PWN σ << 1
terminal shock radius
ΔΩ < 4π if anisotropic windadiabatic losses
PWN radiationsynch. dominates lossesin the wind: Lsyn wind ≈ 0in the PWN Lsyn pwn ∝ nB²Lsyn∝ r2 for RTS ≤ r ≤ Req, Lsyn ∝ r-4 for R > Req ↓↓↓
extTS
wramthpwn p
cREpp =π
== + 2224
&
SNR
PWN
wind
RTSRext
Requipartition
ppwn≈ pSNR
TeVhe
TeV
XIR
CMBhe
XradioBe
syn +γ→ν+
γ<<→ν+
−→+
±
±
±rEEcv:Rr
Erv:Rr
)vr(drdEvEE
eadequip
adequip
ew
ead
32
033
2
22
−=⇒σ=>>
=⇒∝<
−=⋅∇−=
−
&
&
rr&
21 /
ext
wTS cp
ER ⎟⎟⎠
⎞⎜⎜⎝
⎛ΔΩ
=⇒&
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RTS ∼ 2 109 RLC = 0.1 pc (10”) opt. polarization ⇒ Bφ
σ ∼ 0.001 - 0.003 - 0.005from expansion velocity from syn. brightness and Ebreakbetween radio and opticalfrom IC emission
Crab nebulaCrab nebula
cut-off 25 MeVinflection 150 MeVB = 30 nT B(equip, opt) = 30 nTsynchrotron + IC (80% synchro-self-Compton, 20 % IC sur CMB, mm, FIR)
de Jager et al. ’96 Aharonian ‘04
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continuous injection !if syn. losses dominate
synchrotron break depends only on geometry and acceleration efficiency!B(Crab) ~30 nT => Emax ~ 1015 eV => synchrotron lifetime ~ 30 days
⎟⎠⎞
⎜⎝⎛
θα=θ∝ν
⎟⎠⎞
⎜⎝⎛ α⎟
⎠
⎞⎜⎝
⎛ θ=⇒τ≈τ
⎟⎠⎞
⎜⎝⎛ θ
⎟⎠⎞
⎜⎝⎛≈≈τ
πα=τα≈τ⎟
⎠⎞
⎜⎝⎛χα≈τ
−−
−
−−
sinMeVsinEBh rad. synch.
TsinBeV.E
TsinB
eVEs.EE
eBcE
cu)(
max
//
maxsynacc
synsyn
LarmorLarmorvent
v,Bacc C
DD 25
11010146
1110642
2
2
2121
714
216
2
2
&
rr
De Jager et al. ‘96
Crab nebulaCrab nebula
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Crab nebulaCrab nebula
how many pairs?if single Ee
-2.3 power-law for the syn. rad. above the IR => γ1w ~ 2 106
multiplicity ~ ok for pulsar modelsbut it does not explain the radio data thatrequires ~ 1040 e± s-1 over the 952 yr lifetimeif broken power-law (Ee
-1.5 for the radio, Ee-2.3
for opt. to X) ⇒ γ1w ~ 103 and
what acceleration process?Ee
-2.3 ok from relativistic Fermiwhat at lower energies?why a break between the radio and higher ν?
GJ
ewpsrpsr
w
Ns.N
cmNEE
E
&&
&&&
&
4138
21
1010821
≈=⇒
γ=≈σ+
=
−±
±
GJNs.N && 7141 1021075 ≈= −±
VLA
Spitzer
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synchrotron lifetimemaximum distance reached before e± radiate all their energy (roughly)
synchrotron break energy depends on angular distance
PWN size in the Crabradio > opt > X
spectral agingspectral aging 12
11855
−−
⎟⎠⎞
⎜⎝⎛
⎟⎟⎠
⎞⎜⎜⎝
⎛ ><=
TeVE
nTBkyr.tsyn
0TSwTS
maxTS0
TSw
synsyn
TSwwsyn
EE if R
rBYR
c)r(E E(r) to )(RE from integrate
Rr
Yc)r(v withdr
)r(vE
drdrdtEdE
RrYB)r(B withEaBE
<<⎟⎟⎠
⎞⎜⎜⎝
⎛∝⇒
⎟⎟⎠
⎞⎜⎜⎝
⎛===
⎟⎟⎠
⎞⎜⎜⎝
⎛=−=
−
−
5
21
3
2
2
1222
&&
&
9
31
52
2
2
−
⎟⎟⎠
⎞⎜⎜⎝
⎛∝ν⇒
∝ν
TSwTSmax
maxmax
Rr
BYRch
BEh
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spectral agingspectral aging
e± aging with distance=> spectrum softens with r
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spectral agingspectral aging
difficulty: integrated history in the tail when the pulsar has substantially moved
ex: W 44 and PSR B1853+010.5 - 4 keV
4 - 9.5 keV ASCA
8.4 GHz1.442 GHz
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PWN evolutionPWN evolution
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phase I: free spherical expansionphase I: free spherical expansion
during the first hundred yearspulsar motion <<Ėpsr ~ cte ESN = ½ Mej vej
2
pressure-driven expansion since ppwn >> pejecta
adiabatic losses dominate=>
supersonic expansion with increasing speed vpwn ∝ t1/5
PWN isobaric (p = 1/3 ρc2 ⇒ cs=c/√3) ⇒ spherical symmetry
drives a shock in the ejecta
ex: G21.5-0.9 + PSR J1833-1034P + 61.8 ms, age < 1 kyr ≠ τpsr
Ėpsr = 3 1030 W, D ~ 4.7 ± 0.4 kpcRpwn ~ 1 pc, RSNR ~ 3.4 pc
56
3
21103
44
51
31 1010101011
//ej
/SN
/psr
pwnyr
tM
M
JE
W
Epc.R ⎟
⎟⎠
⎞⎜⎜⎝
⎛⎟⎟⎠
⎞⎜⎜⎝
⎛⎟⎠
⎞⎜⎝
⎛⎟⎟⎠
⎞⎜⎜⎝
⎛≈
−
Θ
&
Chevalier ’77 & ‘82
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Interstellar medium
SNR structureSNR structure
ejecta
ISM
shocked ISM
shocked ejecta
contact discontinuity
pec= pMISc
forward shock
reverse shock
reverse reverse shockshock
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Phase II: crushed by SNR reverse shockPhase II: crushed by SNR reverse shock
compression by the SNR reverse shockPWN had freely expanded ⇒ ppwn < pSNR ⇒ large compression
ex: Vela: pulsar not at PWN center nor moving away from it
Blondin ’01
Bucciantini ‘03
Vela 2.4 GHz
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crushed by SNR reverse shockcrushed by SNR reverse shock
asymmetrically crushed if collimated windor ISM anisotropy
High ISMdensity
Low ISMdensity
Blondinet al.
Blondin ‘01
20 kyr 30 kyr 56 kyr
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G0.9+0.1G0.9+0.1
crushed PWN or X-ray hardening from Doppler boosting of jet/torus ?Porquet ’03Gaensler ’02
distance uncertainpulsar not discovered
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G0.9+0.1G0.9+0.1radio 1.5 GHz contours HESS
radio 0.3 GHz contours HESS
G0.9+0.1 SgrA *
TeV(*)ISRFCMBe
X)nT.(Be
TeV
TeV
→++
→+±
±− 5010010
assumingstarlight fieldin central Galaxy
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phase III: expansion in a Sedov SNRphase III: expansion in a Sedov SNR
subsonic expansion in hot ejecta (shocked by RS)SNR in Sedov phase
RFS ∝ t2/5, vFS ∝ t-3/5, pshocked ISM = pshocked ejecta ∝ vFS2 ∝ t-6/5
PWN in pressure equilibrium with surrounding ejectappwn = pshocked ejecta and ppwn ∝ V-γ
ad ∝ V-4/3 ∝ R-4 ⇒ Rpwn ∝ t3/10
deceleration of PWN
simulated example:free expansion t < 1.8 kyrcrushed: 1.8-15 kyr(reverberations)expansion in Sedov SNR :15-30 kyr
Bucciantini ’03
contact discontinuity
103
1511
/pwnpsrini
/pwnpsrini
tR)t(Et
tRcteEt
∝↓⇒τ>
∝=⇒τ<
&
& van der Swaluw ’01
Reynolds & Chevalier ‘84
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phase IV: bow shock confinementphase IV: bow shock confinement
crossing time for the pulsar to cross the SNR shell
Mach number at crossing
pulsar motion becomes supersonic at ½ tcross or 0.68 RSNR
PWN confined by the pressure behind the bow shock
3531
36
31
44
5251
500101044
151
/psr
/ISM
/SN
cross
//
ISM
SNRSedovSNRpsrpsr
s/kmv
mn
JEkyrt
tE.RtvR
−−
− ⎟⎟⎠
⎞⎜⎜⎝
⎛⎟⎠
⎞⎜⎝
⎛⎟⎠
⎞⎜⎝
⎛=⇒
⎟⎟⎠
⎞⎜⎜⎝
⎛ρ
≈==
1335
7443
47
43
52
52
52
212 .
)/p(vM vp
vvvv 4 X for vv
vt
Rt
Rv
/ISM.shISM.sh
relpsrISMISM.shSNRISMISMshth
SNRISM.shpsrrelSNRISM.sh
psrpsrSNR
SNR
==ργ
=⇒ρ=ρρ=
=−=⇒==
===
van der Swaluw ‘03
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forward distance to TS
backward distance to TS
)ISM,,(R
RR
RR
wTS
TSYY
CD
CDXX
BS
Γσ
=
=
−
−
1
1
11
−
±
∝σ⇒
⇒σγ
r)E,(B
)e(N,E,
psrw
psrw
&
&&
TSCDBSY(σ) ratio
X ratio
shocked ISM
shocked windisobaric
σ ≤ 0.1 γw ≥ 103bow shock PWNbow shock PWN
cRE)(pvp
TS
wTSpsrISMBS 23
132
22
24π
+≈=ρ=&
21
24
/
psrISM
wTS
vcER ⎟
⎟
⎠
⎞
⎜⎜
⎝
⎛
ρπ=⇒
&
21
2 4
/
ISM
wTSTSISM cp
ERpp ⎟⎟⎠
⎞⎜⎜⎝
⎛π
=⇒≈&
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bow shock PWNbow shock PWN
• analytical profile• cooling thin shell approximation, ok if M >> 1• Wilkin ‘96
⎟⎟⎠
⎞⎜⎜⎝
⎛θθ
−θ
=θtgsin
R)(R BSBS 130
17
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magnetotailmagnetotail
toroidal field and reconnection
Romanova et al. ‘05
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bow shock PWN in SNRbow shock PWN in SNRCXO 0.3-10 keV
radio 8.5 Ghz
IC 443
RCD ~ 0.06 pc
XMM 0.2-12 keV radio 0.8 Ghz
G327.1-1.1
Gaensler & Slane ‘06
Gaensler et al. ‘06
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breaking through the shellbreaking through the shell
PWNe survive the crossing of the SNR shell
ex: the BoomerangG106.3+2.7 + PSR J2229+6114
Ė = 2.25 1030 W (2nd to Crab)D = 0.8-3 kpcSNR = 14 × 6 pc at 800 pcPWN 0.8 pc long breaking through?bizarre SNR…
van der Swaluw ‘03
Kothes et al. ‘01
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phase V: bow shock PWN in ISMphase V: bow shock PWN in ISM
M >> 1 in ISM (cs = 1-10 km/s)ex: the Mouse
G359.23-0.82 + PSR J1747-2958P = 99 ms, 25.5 kyr, Ėpsr = 2.5 1029 W, D ~ 5 kpcRCD ~ 0.02 D5 pc
radio from highlycompressed B behindthe front TS (severe losses=> only radio emitting e±
survive)radio + X rays fromweakly compressed Bbehind the rear TS
as in W44
Gaensler et al. ‘04
19
I. Grenier
Guitar PSR 2224+65Guitar PSR 2224+65
the « breathing » PWNHα bow shock from recombining HII
photoionized gas and charge exchangevariability over 7 years, not from the wind
but loosened confinement (n ↓) ⇒ decreased head flux
Chatterjee & Cordes ’04
I. Grenier
inner wind asymmetriesinner wind asymmetries
jets and torii
20
I. Grenier
jet 0.25 pcterminal shock:
ΔΩ ~ 1 sr => RTS ~ 0.4 pc (40’’)ΔΩ ~ 4π sr => RTS = 0.1 pc (10’’)wisps à 8’’ et 12’’
Doppler amplification & shiftwisp waves travel at 0.5 c
Crab torus and jetsCrab torus and jets
WL
W.
pwn30
31
102
10654
=
=ΩΩ &I
CXO HST November-April
I. Grenier
3C 58 (SN 1181)3C 58 (SN 1181)65.7 ms τpsr = 0.8 kyr ≠ τkin = 5.4 kyr CXO X-ray colours
21
I. Grenier
Vela et PSR 1509-58Vela et PSR 1509-58
CXO
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Contours: RosatGreyscale: Radio
PSR B1509-58 MSH 15-52PSR B1509-58 MSH 15-52
1.7 kyr old pulsar at 5.2 ± 1.4 kpc, Ėpsr = 1.8 1030 WX-ray jets (12 pc, 0.5 c, long-term variability) and variable knots (0.6c)2 TeV arms ± 25 pc long, radiatively more efficient than the Crab
DeLaney ’06Aharonian ‘05
22
I. Grenier
Photon index2.27 ± 0.03 ± 0.20
MSH 15-52MSH 15-52
TeV)IR(ISRFCMBe
X)nT.(Be
TeV
TeV
→++
→+±
±− 7110010
I. Grenier
high-energy relic tailshigh-energy relic tails
23
I. Grenier
Vela tailVela tail
~ 5 pc long tailseen up to 50 TeVE-1.9 ou E-1.5+cutoff
Vela pulsar
ROSATcontours
Preliminary
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PSR B1823-13 - G18.0-0.7PSR B1823-13 - G18.0-0.721.4 kyr old pulsar
P = 101 ms, Ėpsr = 2.8 1029 W, D = 3.9 ± 0.4 kpcTeV tail LTeV ~ 50 pc: e±(TeV) + CMB-IR → γ(≥ 200 GeV)X-ray tail LX ~ 5 pc: e±(200 TeV) + Bequip(1 nT) → X
LTeV/LX ok in uniform Bbut very high vpsr or v± required
HESS
Aharonian et al. ’05 Gaensler et al. ‘03
okEE
)E()E(
LL
eTeV
eX
eXsyn
eTeVsyn
X
TeV 200≈=τ
τ<
24
I. Grenier
0.02
pc
0.1 pcTeV E
smeVE.F
E.L
E.L
planesky the to i
pc D,kyr,W.E
eX
..eVX
psrkeV.
psrsr,GeV.
psr
100
62
1091
1061
30
1603401033
122040
6530
21210
27
≈
=ν
=
=
°≤
==τ=
−−±+ν
−−
−−
&
&
&
Caraveo ’03the wings of Gemingathe wings of Geminga
pc.Rpc .Rpc .R
BS
CD
TS
030020010
=
=
=
11
−
±
∝σ⇒
⇒σγ
r)E,(B
)e(N,E,
psrw
psrw
&
&&
TSCDBSBSBSY=2.3
X=3.1
1.9 nT2.1
nT
105 m-3, 8000 K0.6 nT, 135 km/s
Forot & Grenier ’06
I. Grenier
XMM image => σ << 1acceleration site?
Fermi at BS:
Fermi at TSEmax = e ΔΦopen = 400 TeV
emissivity problem
Distance from the pulsar (0.01pc)-5 -4 -3 -2 -1 0 1 2 3 4 5
Dis
tanc
e fr
om th
e pu
lsar
(0.0
1pc)
-3
-2
-1
0
1
2
3
3410
3510
3610
Photons
Distance from the pulsar (0.01pc)-5 -4 -3 -2 -1 0 1 2 3 4 5
Dis
tanc
e fr
om th
e pu
lsar
(0.0
1pc)
-3
-2
-1
0
1
2
3
3410
3510
3610
Photons
Distance from the pulsar (0.01pc)-5 -4 -3 -2 -1 0 1 2 3 4 5
Dis
tanc
e fr
om th
e pu
lsar
(0.0
1pc)
-3
-2
-1
0
1
2
3
3410
3510
3610
Photons
σ = 1 σ = 0.1σ = 0.01
i = 0°the wings of Gemingathe wings of Geminga
GeVERRu
D but ok E maxCDBSISM
BohmXX
301
112
<⇒−≤≈−+
−
TeVE but ok E max. 3522 <−
/amplifiedcompressed B unless LL ),E(N XobsXpsr <<⇒σ±&&
25
I. Grenier
Kookaburra and RabbitKookaburra and Rabbit HESS
1.5 GHz