1

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Pulsar wind nebulaePulsar wind nebulae

Isabelle A. GrenierUniversité Paris 7 & CEA Saclay

I. Grenier

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

16

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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

18

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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

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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

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inner wind asymmetriesinner wind asymmetries

jets and torii

20

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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

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3C 58 (SN 1181)3C 58 (SN 1181)65.7 ms τpsr = 0.8 kyr ≠ τkin = 5.4 kyr CXO X-ray colours

21

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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

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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