rino bandiera, oaafundamental physics & astrophysics of snrssna07, may 20-26, 2007 fund. physics...

Rino Bandiera, OAA Fundamental Physics & Astrophysics of SNRs SNA07, May 20-26, 2007

Fund. Physics & Astrophysics

of Supernova Remnants• Lecture #1

– What SNRs are and how are they observed– Hydrodynamic evolution on shell-type

SNRs– Microphysics in SNRs – electron-ion equ

• Lecture #2– Microphysics in SNRs - shock acceleration– Statistical issues about SNRs

• Lecture #3– Pulsar wind nebulae


Order-of magn. estimates• SN explosion

– Mechanical energy:

– Ejected mass:

• VELOCITY:

• Ambient medium– Density: Mej~Mswept when:

• SIZE:

• AGE:

erg1051SN E

Sun33

ej 5g10 MM

19ejSNej scm10/ MEV

3ISM cm1.0 n

pc3cm104/3 193/1ISMejSNR nMR

yr30010/ 10ejSNRSNR sVRt


“Classical” Radio SNRs• Spectacular shell-like morphologies

– comparedto optical

– polarization– spectral index

(~ – 0.5)

BUT

• Poor diagnostics on the physics– featureless spectra (synchrotron emission)– acceleration efficiencies ?

Tycho – SN 1572


90cm Survey 4.5 < l < 22.0 deg (35 new SNRs found; Brogan et al. 2006)

Blue: VLA 90cm Green: Bonn 11cm Red: MSX 8 m

• Radio traces both thermal and non-thermal emission

• Mid-infrared traces primarily warm thermal dust emission

A view of Galactic Plane


SNRs in the X-ray window• Probably the “best”

spectral range to observe

– Thermal:• measurement of

ambient density

– Non-Thermal:• synchrotron-emitting

electrons are near the maximum energy (synchrotron cutoff)

keV12ejee VmkT

dVnnEM eH


X-ray spectral analysis• Low-res data

– Overall fit with thermal models

• High-res data– Abundances of

elements– Single-line

spectroscopy!


Shell-type SNR evolutiona “classical” (and wrong) scenario

Isotropic explosion and further evolutionHomogeneous ambient medium

Three phases:• Linear expansion• Adiabatic expansion• Radiative expansion

Isotropic

Homogeneous

Linear

AdiabaticRadiative


Basic concepts of shocks• Hydrodynamic (MHD)

discontinuities• Quantities conserved

across the shock– Mass– Momentum– Energy– Entropy

• Jump conditions(Rankine-Hugoniot)

• Independent of the detailed physics

12

1122

22 pVpV 12

11122

222 2/2/ wVVwVV

1122 VV

12 ss

shock111 V,,p222 V,,p

V

4/3;4/;4 21121212 VpVV

If 3/5

2111 Vp Strong shock

21121212 1

2;

1

1;

1

1VpVV


Den

sity

Radius

Forward shock

Reverse shock

Forward and reverse shocks

• Forward Shock: into the CSM/ISM (fast)• Reverse Shock: into the Ejecta (slow)


Dimensional analysisand Self-similar models

• Dimensionality of a quantity:• Dimensional constants of a problem

– If only two, such that M can be eliminated, THEN evolution law follows immediately!

• Reduced, dimensionless diff. equations– Partial differential equations (in r and t)

then transform into total differential equations (in a self-similar coordinate).

rqp TLMA


Early evolution• Linear expansion only if ejecta

behave as a “piston”• Ejecta with and• Ambient medium

with and • Dimensional parameters

and• Expansion law:

trV / ntrtg )/(3ej

0V sqr amb

)3()3( nn TMLg )3( sMLq

)/()3()/(1/ snnsn tqgR


A self-similar model• Deviations from

“linear” expansion

• Radial profiles– Ambient medium– Forward shock– Contact

discontinuity– Reverse shock– Expanding ejecta

60.0:7,2 tRns 90.0:12,2 tRns

(Chevalier 1982)


Evidence from SNe• VLBI mapping (SN 1993J)

• Decelerated shock

• For an r -2 ambient profileejecta profile is derived


The Sedov-Taylor solution• After the reverse shock has reached

the center• Middle-age SNRs

– swept-up mass >> mass of ejecta– radiative losses are negligible

• Dimensional parameters of the problem

• Evolution:• Self-similar, analytic solution (Sedov,1959)

3ISMISM : ML 22

SNSN : TMLEE

5/25/1ISMSNSNR )/()( tEtR


The Sedov profiles

• Most of the mass is confined in a “thin” shell• Kinetic energy is also confined in that shell• Most of the internal energy in the “cavity”

Density

Temperature

Pressure


Thin-layer approximation• Layer thickness

• Total energy

• Dynamics

1233

44

2

11

32

2 RRrRrR

2c13

22c3 ;

3

4;

213

4ppRM

uM

pRE

223c

22 3

14 RRRR

dt

dpRMu

dt

d

2

1

5

2;

3

14

q

q

qtR q

2

5

15

22

13 1

1

2

)1)(1(

2

5

2

2

1

3

4

t

R

t

RRE

12.1

3

5

Correct value: 1.15 !!!


What can be measured (X-rays)

pc5.12 5/24

5/10

5/151Sed tnER

dVnnEM eH shockx 28.1 TT

from spectral fits

d

t

n

E

VkT

dR

dEM

x

0

2

/

/

… if in the Sedov phase


SN 1006 Dec.Par. = 0.34Tycho SNR (SN 1572) Dec.Par. = 0.47

Testing Sedov expansion

Required:• RSNR/D (angular size)

• t (reliable only for historical SNRs)

• Vexp/D (expansion rate, measurable only in young SNRs)

5/2/ SNRexp RtV

Deceleration parameter


Other ways to “measure”the shock speed

• Radial velocities from high-res spectra(in optical, but now feasible also in X-rays)

• Electron temperature from modelling the (thermal) X-ray spectrum

• Modelling the Balmer line profile in non-radiative shocks (see below)


End of the Sedov phase

• Sedov in numbers:

• When forward shock becomes radiative: with

• Numerically:

117/20

17/15117/7

017/5

51tr

17/90

17/451

4tr skm260

pc19

yr109.2

nEV

nER

nEt

0coolagetr

1:

nttt

pc5.12 5/24

5/10

5/151Sed tnER

13116 scmerg10)( TT


Beyond the Sedov phase• When t>ttr, energy no longer conserved.

What is left?• “Momentum-conserving

snowplow” (Oort 1951)

• WRONG !! Rarefied gas in the inner regions

• “Pressure-driven snowplow” (McKee & Ostriker 1977)

4/13

tRconst

constVR

ISM

ISM

Kinetic energy

Internal energy)33/(2

2ISM

3kin

3inninn

3int

/

/

tR

VRE

RPRE

3/5for7/2 tR


Numerical results

ttr

Blondin et al 1998

2/5 0.33

2/7=0.29

1/4=0.25

(Blondin et al 1998)


An analytic model• Thin shell approximation

• Analytic solution

R

Rp

td

pdRp

td

RMdRR

td

Md

cc2

c2

0 3;4)(

;4

13

2

)2(33

KRRRR

632 HRRKR

H either positive (fast branch)limit case: Oort

or negative (slow branch)limit case: McKee & Ostriker

H, K from initial conditions

Bandiera & Petruk 2004


Inhomogenous ambient medium

• Circumstellar bubble (ρ ~ r -2)– evacuated region around the star– SNR may look older than it really is

• Large-scale inhomogeneities– ISM density gradients

• Small-scale inhomogeneities– Quasi-stationary clumps (in optical) in

young SNRs (engulfed by secondary shocks)

– Thermal filled-center SNRs as possibly due to the presence of a clumpy medium


Collisionless shocks• Coulomb mean free path

– Collisional scale length (order of parsecs)– Larmor radius is much smaller (order of km)

• High Mach numbers– Mach number of order of 100

• MHD Shocks– B in the range 10-100 μG

• Complex related microphysics– Electron-ion temperature equilibration– Diffusive particle acceleration– Magnetic field turbulent amplification


Electron & Ion equilibration• Naif prediction, for collisionless shocks

• But plasma turbulence may lead electrons and ion to near-equilibrium conditions

• Coulomb equilibration on much longer scales

ii

ee

i

i

e

e Tm

mTV

m

kT

m

kT 2

sh16

3

(Cargill and Papadopoulos 1988)

ii

ee T

m

mT

pcskm1000cm1

4.24

1sh

1

30

2/5

Vn

T

TL

p

eeq

c.g.s.13.02/3

e

epe

e

T

TTn

dt

dT

(Spitzer 1978)


Optical emission in SN1006• “Pure Balmer” emission

in SN 1006

• Here metal lines are missing (while they dominate in recombination spectra)– Extremely metal deficient ?


“Non-radiative” emission• Emission from a radiative shock:

– Plasma is heated and strongly ionized– Then it efficiently cools and recombines– Lines from ions at various ionization levels

• In a “non-radiative” shock:– Cooling times much longer than SNR age– Once a species is ionized, recombination is

a very slow process

• WHY BALMER LINES ARE PRESENT ?


The role of neutral H• Scenario: shock in a partially neutral gas• Neutrals, not affected by the magnetic

field, freely enter the downstream region• Neutrals are subject to:

– Ionization (rad + coll) [LOST]– Excitation (rad + coll) Balmer narrow– Charge exchange (in excited lev.)Balmer broad

(Chevalier & Raymond 1978, Chevalier, Kirshner and Raymond 1980)

•Charge-exchange cross section is larger at lower vrel

•Fast neutral component more prominent in slower shocks


H-alpha profiles

(Hester, Raymond and Blair 1994)

(Kirshner, Winkler and Chevalier 1987)

Cygnus Loop

•FWHM of broad component (Ti !!)

•FWHM of narrow component

• (T 40,000 K – why not fully ionized?)

MEASURABLE QUANTITIES

•Intensity ratio

•Displacement (not if edge-on)


SNR 1E 0102.2-7219• Very young and bright SNR in the SMC• Expansion velocity (6000 km s-1, if linear expansion)

measured in optical (OIII spectra) and inX-rays (proper motions)

• Electron temperature~ 0.4-1.0 keV, whileexpected ion T ~ 45 keV

• Very small Te/Ti, or Ti

much less than expected?Missing energy in CRs?

(Hughes et al 2000, Gaetz et al 2000)

Optical

Radio

X-rays


Lectures #2 & #3• Shock acceleration

– The prototype: SN 1006– Physics of shock acceleration– Efficient acceleration and modified shocks

• Pulsar Wind Nebulae– The prototype: the Crab Nebula– Models of Pulsar Wind Nebulae– Morphology of PWN in theory and in

practice– A tribute to ALMA


The “strange case” of SN1006

Tycho with ASCA

Hwang et al 1998

“Standard”X-ray spectrum


Thermal & non-thermal• Power-law spectrum at the rims• Thermal spectrum in the interior


shock

X

flow speed

(in the shock reference frame)

Diffusive shock acceleration

• Fermi acceleration– Converging flows– Particle diffusion

(How possible, in acollisionless plasma?)

• Particle momentum distributionwhere r is the compression ratio (s=2, if r = 4)

• Synchrotron spectrum• For r = 4, power-law index of -0.5• Irrespectively of diffusion coefficient

srr pppF )1/()2()(

)1/(2/32/)1()( rsS


The diffusion coefficient• Diffusion mean free path

(magnetic turbulence)

(η > 1)

• Diffusion coefficient

eB

mcr

2

g 2

resg )/(with BBr

eB

mcv

33

3

fu

x

fF diffdiff


…and its effects• Acceleration time

• Maximum energy

• Cut-off frequency

– Naturally located near the X-ray range– Independent of B

2sh

3

2121acc )1(

)1(113

eBu

mc

r

rr

uuuut

BuB

tBu

t sh

/

1 2max2syn2

shacc

2sh2

maxcutoff

uB

keVskm10

11.02

13sh

cutoff

u


Basics of synchrotron emission

• Emitted power • Characteristic frequency • Power-law particle distribution• If then• Synchrotron life time

221

222

32

4

syn )(sin3

2 BcmcBcm

eW

22

2syn 2

sin3

2

29.0

Bcmc

eB

sppF )( 2/)1()( sS

dt

dmcW

)( 2

syn

Bct

1syn

1


SN 1006 spectrum• Rather standard ( -0.6) power-law

spectrum in radio(-0.5 for a classical strong shock)

• Synchrotron X-rays below radio extrapolation

Common effect in SNRs (Reynolds and Keohane 1999)

• Electron energy distribution:

• Fit power-law + cutoff to spectrum:

“Rolloff frequency”

)/exp()( maxEEEEN se

))/(exp()( 2/1rolloff S


Measures of rolloff frequency

• SN 1006 (Rothenflug et al 2004)

• Azimuthal depencence of the break

Changes in tacc? or in tsyn? η of order of unity?


Dependence on B orientation?

• Highly regular structure of SN 1006.Barrel-like shape suggested (Reynolds 1998)

• Brighter where B is perpendicular to the shock velocity?

Direction of B ?


Radio – X-ray comparison

•Similar pattern (both synchrotron)

•Much sharper limb in X-rays (synchrotron losses)

(Rothenflug et al 2004)


(Rothenflug et al 2004)

• Evidence for synchrotron losses of X-ray emitting electrons

• X-ray radial profile INCONSISTENT with barrel-shaped geometry (too faint at the center)


3-D Geometry. Polar Caps?

Ordered magnetic field

(from radio polarization)

Polar cap geometry:electrons acceleratedin regions with quasi-parallel field

(as expected from the theory)


Statistical analysis(Fulbright & Reynolds 1990)

Barrel-like SNR(under variousorientations)

Polar cap SNR(under variousorientations)

Expected morphologies in radio


The strength of B ?

• Difficult to directly evaluate the value of the B in the acceleration zone.

νrolloff is independent of it !

• “Measurements” of B must rely on some model or assumption


Very sharp limbs in SN 1006

ASCA

Chandra


B from limb sharpness

Profiles of resolved non-thermal X-ray

filaments in the NE shell of SN 1006

(Bamba et al 2004)

Length scales 1” (0.01 pc) upstream 20” (0.19 pc) downstream

Consistent withB ~ 30 μG


A diagnostic diagram• Acceleration time

tacc = 270 yr• Derivation of

the diffusioncoefficients:u=8.9 1024 cm2s-1

d=4.2 1025 cm2s-1 (Us=2900 km s-1)to compare withBohm=(Emaxc/eB)/3

rolloff

tsync> tacc

> Bohm


Non-linear shock acceleration

• Such high values of B are not expected in the case of pure field compression(3-6 μG in the ISM, 10-20 μG in the shock – or even no compression in parallel shocks)

• Turbulent amplification of the field?• Possible in the case of efficient shock

acceleration scenario: particles, streaming upstream, excite turbulence

(e.g. Berezhko; Ellison; Blasi)


Shock modificationDynamical effects of the

accelerated particles ontothe shock structure

(Drury and Voelk 1981)

•Intrinsically non linear

•Shock precursor

•Discontinuity (subshock)

•Larger overall compression factor

•Accelerated particle distribution is no longer a power-law


Deviations from Power-Law• In modified shocks,

acc. particles withdifferent energiessee different shockcompression factors.Higher energy Longer mean free path Larger compress.factor Harder spectrum

• Concavity in particledistribution.

(also for electrons)

Standard PL

Thermal

Blasi Solution


Gamma-ray emission• Measurement of gamma-ray

emission, produced by the same electrons that emit X-ray synchrotron, would allow one to determine the value of B.

SynchrotronIC

Radio X-ray γ-ray

νFν


• On the other hand, there is another mechanism giving Gamma-ray emission– accelerated ions– p-p collisions– pion production– pion decay (gamma)

• Lower limit for B

• Need for “targets”

(molecular cloud?) • Efficiency in in accelerating ions?

(The origin of Cosmic rays)

(Ellison et al 2000)


TeV telescopes generation• H.E.S.S. Cherenkov telescopes

• Observations :• RX J0852.0-4622 (Aharonian et al 2005) • Upper limits on SN 1006 (Aharonian et al 2005)

• RX J1713.7-3946 (Aharonian et al 2006)


Observ. of RX J0852.0-4622

•Good matching between X-rays and gamma-rays

•CO observation show the existence of a molecular cloud

•Pion-decay scenario slightly favoured. Nothing proved as yet


Indirect tests on the CRs• Some “model-dependent” side effects of efficient

particle acceleration• Forward and reverse shock are closer, as effect of

the energy sink• HD instabilities behavior depends on the value of eff

(Decourchelle et al 2000)

(Blondin and Ellison 2001)


Shock acceleration efficiency• Theory predicts (~ high) values of the

efficiency of shock acceleration of ions.• Little is known for electrons• Main uncertainty is about the injection

process for electrons– Shock thickness determined by the mfp of

ions (scattering on magnetic turbulence)– Electrons, if with lower T, have shorter mfps– Therefore for them more difficult to be

injected into the acceleration process


The Σ–D Relationship

• Empirical relation– SNR surface

brightness, in radio– SNR diameter– Any physical

reason forthis relation ?

(Case & Bhattacharya 1998)

A Caswith 64.2

A Caswithout 38.2


A basic question• Is the correlation

representative of the evolution of a “typical object”?

• Or is, instead, the convolution of the evolution of many different objects?

• Theorists attempts to reproduce it.

Berezhko & Voelk 2004


Dependence on ambient density

• Primary correlations are D-n, and Σ-n

• Diff. ISM conditions

(Berkhuijsen 1986)


Crab Nebula – Hcont

Crab Nebula - radio

X-rays

The “Prototype”The Crab Nebula• Optical

Thermal filamentsAmorphous compon.

• RadioFilled-center nebulaNo signs of shell

• X-raysMore compact neb.Jet-torus structure


The Crab Nebula spectrum

-0.3

-1.1-0.8

-1.5

Synchrotron emission

)mG 3.0( neb B

(apart from optical filaments and IR bump)

• Radio

•Optical

•Soft X-rays

•Hard X-rays


Some basic points

• Synchrotron efficiency– 10-20% of pulsar spin-down power

• Powered by the pulsar• High polarizations (ordered field)• No signs of any associated shell.


Basics of synchrotron emission

• Emitted power • Characteristic frequency • Power-law particle distribution• If then• Synchrotron life time

221

222

32

4

syn )(sin3

2 BcmcBcm

eW

22

2syn 2

sin3

2

29.0

Bcmc

eB

sppF )( 2/)1()( sS

dt

dmcW

)( 2

syn

Bct

1syn

1


Simple modelling• Homogeneous models (no info on

structure)• Magnetic field evolution

– Early phases (constant pulsar input)

– Later phases (most energy released)

(Pacini & Salvati 1973)

13

32

26; t

R

LtB

tLRBWt B

22

2 )()()(6; t

R

RLBRWBRt B


• Power-law injection– With upper energy cutoff– Continuum injection

• link to the pulsar spin down

• Particle evolution (adiabatic vs synchrotron losses)• Evolutionary break

• Adiabatic regime(-0.3 in radio)

• Synchrotron-dominated regime(-0.8 in optical)

2321

22br2br2

1brsyn )(

)()(

1

ttBc

ctBc

ttBctt

5.121 s

sdttjtN ),(),(

1syn),(),( sttjtN

2/)1()( sS

2/)( sS


Kennel & Coroniti model (1984)

Basics of “Pulsar Wind Nebula” scenario• Pulsar magnetosphere• Pulsar wind• Termination shock• Pulsar Wind Nebula• Interface with the

ejecta (CD, FS)• Stellar ejecta• Interface with the

ambient medium(RS, CD, FS)

• Ambient medium (either ISM or CSM)

Pulsar magnetosphere

Pulsar wind

Termination shockPulsar Wind Nebula

Stellarejecta

ISM


The ingredients• Pulsar wind

– super-relativistic– magnetized

(toroidal field)– isotropic

• Termination shock– mass conservation– magnetic flux cons.– momentum cons.– energy cons.

where (specific enthalpy)

2

2

particle

Poynting

4 mcnu

B

F

F

speed-4 comp radialu

densityproper

n

2211 unun

EuBuB 222111 //

8/8/ 222

2222

211

2111 BpunBpun

4/4/ 2222211111 EBunEBun

n

pmc

12


Large and small σ limits• Large σ

– weak shock– flow stays super-relativistic– neither field, nor density jump– inefficient in converting kinetic into

thermal energy

• Small σ– strong shock– flow braked to mildly relativistic speed– both field and density increase– kinetic energy efficienly converted

8/,//,, 12

21122122 umckTnnBB

18/,/3/,3,8/9 12

21122122 umckTnnBB


MHD evolution in the nebula• Steady solution (flow timescale << SNR age)

– number flux cons. - magnetic flux cons.– momentum cons. - energy cons.

• Asymptotic velocity !!!– no solution for V∞=0

– outer expansion Vext~1500 km s-1 (for the Crab Nebula)

– then σ~3 10-3

– size of termination shock, from balance of wind ram pressure and nebular pressure

Rn~10 arcsec

(wisps region)

1

ucV

c

V

R

R

R

VLR

cR

L ext

n

s3n

extn2s 3/4

/

4


Radial profiles

• Inner part with:• Outer part with:• Equipartition in the outer part:

rBconstnru ,,2

12 ,, rBrnconstu

042

2

4rr

mcnu

B


Do we expect what observed?

• Injected particles– power-law, between a min and a max energy

– only 1 free parameter (n2 and p2 from the jump conditions at the termination shock)

– plus wind parameters (L, σ and γ1 )

• Energy evolution during radial advection

222)12(

22 ;)( Af

2213

Bcdr

dnu

ndr

du


Best-fit solution• Parameters:

• Fit to:

003.0,6.0,103,cm103,serg105 61

17138 srL


Problems -Ia• The sigma paradox

– A value is required, in order to get an effective slowing-down of the flow, and a high (10-20 %) synchrotron conversion efficiency

– BUT the (magnetically driven) pulsar wind cannot have been produced with a low σ .

– With a normal MHD evolution, the value of σ must keep constant from the acceleration region till the termination shock.

1


Problems - Ib• A POSSIBLE WAY OUT

– A tilted pulsar generates a striped wind.


Problems -Ic– Magnetic reconnection in the wind zone

(if possible) would dissipate the field.(Coroniti 1990)

– Reconnection in the wind zone does not efficiently destroy the field. Reconnection at the termination shock is more effective.

(Lyubarski & Kirk 1991)


Problems - IIa• The unexpected radio emission

– Predicted radio flux is far lower (a factor ~100) than observed.

– No easy way to cure it. Little freedom on the particle number. Total power is fixed: more particles mean a lower γ1.

– Radio emitting electrons as a relict. Was the Crab much more powerful in the past? Ad hoc. All PWNe are radio emitters.


Problems IIb– Can it be “Diffusive synchrotron

radiation”?(Fleishman & Bietenholz

2007)

Turbulence spectral index ν.– Theory only for a fully turbulent field

• Total spectrumis reproduced

• But observedpolarization isnot explained


Non-spherical structure

• Particle, moving passively along field lines (flow motion assumed to be irrotational)

• Axisymmetric nebular field structure• Steady state solutions

(Begelman & Li 1992)


van der Swaluw 2003

pulsar axis

3C 58

MHD simulations


Elongated structures of PWNe

G5.4-0.1

Crab Nebula

pulsar spinG11.2-0.3

3C 58


Details of the structure

knot

jet

inner ring

torus

counter-jet

Crab NebulaVela


Crab Nebula (Weisskopf et al 2000)

40” = 0.4 pc

PSR B1509-58 (Gaensler et al 2002)

4’ = 6 pc

3C 58 (Slane et al. 2004)

13” = 0.2 pc80” = 0.8 pc

Vela Pulsar (Pavlov et al. 2003)

Jet sizes


Simulating PWNe

• Relativistic MHD codes• Modelling a PWN like the Crab

Velocity Magnetization Max Energy

(Komissarov, 2006; Del Zanna et al 2004, 2006)


Surface brightness mapsJet-Torus structure


Ingredients• Wind parameters

– magnetization (still small, but not too much)σ~0.02 – 0.1 aaa

– wind anisotropy ( γeq~10 γpol )

– “filling” the jets (since B = 0)


ISM

Sh

ock

ed ISM

Sh

ock

ed E

ject

a

Unsh

ock

ed

Eje

cta

PW

N

Puls

ar

Win

d

Forward Shock

Reverse ShockPWN Shock

PulsarTerminationShock

PWN-ejecta interaction• PWNe are confined by the associated

shell-like SNR• Not only the SNR is detectable (like

in the Crab)• In the Crab Nebula

UV emissionassociated with aslow shock (againstthe SN ejecta)


A TRIBUTE TO ALMA• SNRs and PWNs are mostly non-

thermal in that spectral range.– no use of spectral capabilities– use of high spatial resolution, + wide field,

+ photometric stability (extended sources)

• Is mm-submm a “new band” for SNRs, or just an extension of the radio range?

• A study of the Crab Nebula(extension of a former work, Bandiera et al 2002)


What has been done already

• Comparison of 1.3 mm (230 GHz) images (with IRAM 30-m telescope, 10” res) and radio (20 cm, VLA) maps

230 GHz map Spectral map

-0.28

-0.20


A further emission component

• Radio spectral index: -0.27• Concave spectral index from radio to mm

Real effector artifact?(absolutephotometry)

• Evidence foran additionalemission component


Component B

• Image obtained optimizing the subtraction of amorphous part, and filaments, of radio image (PSF matched), with best-fit weights.


The subtracted components• Amorphous component: consistent with

an extension of the spectrum to mm, with the radio spectral index (-0.27).

• Filaments: consistent with spectral bending (νb~80 GHz).

• Morphologically, component B resembles more the Crab in the optical than in the radio (ALTHOUGH, in the mm range, electrons of Component B do not lose energy significantly by synchrotron).


The integrated spectrum• Radio comp (A)• Component B,

with low freqcutoff.

• Evidence higherthan from theerror bar.

• Components Aand B coexistin the optical.


Physical scenario• Number of particles in Component B:

Ntot ~ 2 1048.

• Consistent with Kennel & Coroniti)• Filament magnetic fields ~6 times

higher than the rest AND particle do not diffuse in/out of filaments (κ<100 κB).


With ALMA• The same analysis, with a resolution 100

times higher.• Detailed mapping of Component B.• Separation of comp A and B also through

differences in the polarization patterns.• Analysis of the spectral bending in

individual filaments, and possibly even across the filament (B estimates).

• Mapping B in filaments (aligned? ordered?)

rino bandiera, oaafundamental physics & astrophysics of snrssna07, may 20-26, 2007 fund. physics...

Documents

rino bandiera

strong shock slide

derived slide

detailed physics shock

ejecta slow slide

sedov phase slide

snrs lecture

new snrs