galactic cosmic rays and the turbulent heliospheric...

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galactic cosmic rays andthe turbulent heliospheric tail

Paolo Desiati1,2 & Alexander Lazarian2

1 WIPAC - Wisconsin IceCube Astrophysics Center2 Department of Astronomy

University of Wisconsin - Madison

Midwest Magnetic Fields Workshop, Madison, WIApril 4th, 2012

Friday, April 6, 2012

Paolo Desiati

cosmic rays spectrum

≈E-2.7

≈E-3.1

≈E-2.7

• spectral structure & mass composition hold information on

‣ origin of cosmic rays

‣ propagation from sources to Earth

‣ anisotropy in arrival distribution

‣ energy dependence

‣ angular scale

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1 GeV - 1 TeV 1 TeV - 1 PeV > 1 PeV extra-galactic


Paolo Desiati

cosmic ray acceleration in supernova remnants

SN1006

W. Baade & F. Zwicky, Physical Review 46, 76, 1934

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• diffusive shock acceleration in galactic SNR (Baade & Zwicky, 1934 & Fermi, 1949)

φCR =cnCR(E)

4π

nCR(E) ≈ E−γ

RSN

2πR2d

· H

D(E) density of cosmic rays

cosmic ray flux

energy emitted by one SN

cosmic ray acceleration efficiency

radius of galactic disk

rate of supernovae in the Galaxy

propagation term

D(E) ∝ Eδ diffusion coefficient

φCR ≈ 2.4 ·�

ESN

1051erg

�

·�CR

·�

15kpc

Rd

�2

·�

RSN

30yr

�

·(γ − 2) · 3−δ

·�

E

1TeV

�−γ−δ

[TeV −1m−2s−1sr−1]

X-ray (Chandra)opticalradio


Paolo Desiati

cosmic rays observationsall-particle spectrum

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p ≈ E-2.66

He ≈ E-2.58

PamelaAdriani et al. (2011)

CREAMAhn et al. (2010)

p ≈ E-2.80

He (×0.1) ≈ E-2.71

p ≈ E-2.67

He ≈ E-2.48


Paolo Desiati

cosmic rays observationsall-particle spectrum

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residual all-particle spectrum

≈E-3.0

≈E-3.1

PamelaAdriani et al. (2011)

p ≈ E-2.80

He (×0.1) ≈ E-2.71

p ≈ E-2.67

He ≈ E-2.48

KASCADE-GrandeArtega-Velàzquez et al. (2010)

Gaisser & StanevPDG


Paolo Desiati

cosmic rays observationsanisotropy

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05101520

α (hr)

24

0〫

60〫δ (deg)

05101520

α (hr)

24

0〫

50〫

δ (deg)

05101520

α (hr)

24

0〫

50〫

-50〫

δ (deg)

05101520

α (hr)

24

0〫

60〫

δ (deg)

equatorial coordinates

Nagashima et al. (1998)Hall et al. (1999)

ARGO-YBJZhang et al. (2009)

Tibet ASγAmenomori et al. (2006)

Super KamiokandeGuillian et al. (2007)

MilagroAbdo et al. (2009)

< 100 GeV 4 TeV

10 TeV

4 TeV

5 TeV

-90˚

α (hr)0510152024

IceCubeAbbasi et al. (2010)

20 TeV


Paolo Desiati


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equatorial coordinates relative intensity

Amenomori et al. (2011)

Abbasi et al. (2012)

Tibet-ASγ 5 TeV

IceCube-59 20 TeV


Paolo Desiati

anisotropy vs. energy

• CR anisotropy changes phase ~100 TeV

• global amplitude is modulated

Tibet ASγ Amenomori et al., ApJ. 626, L29, 2005

IceCube-22 Abbasi et al., ApJ, 718, L194, 2010

EAS-TOP Aglietta et al., ApJ, 692, L130, 2009

IceCube-59 Abbasi et al., ApJ, 746, 33, 2012

ARGO-YBJ Zhang 31st ICRC Łódź-Poland,2009

360˚

-90˚

0˚

-90˚

360˚ 0˚

20 TeV

400 TeV

IceCubeAbbasi et al. (2012)

δfluctuations =3

23/2

1π1/2

D(E)Hc

D(E) ∝ Eδ

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Paolo Desiati

anisotropy vs. angular scale

detector acceptance

direct exp measurement+ detector acceptance

-0.2

-0.18

-0.16

-0.14

-0.12

-0.1

-0.08

-0.06

-0.04

-0.02

0

0.02

0.04

0.06

0.08

0.1

x 10-2

0 20 40 60 80 100 120 140 160 180 200 220 240 260 280 300 320 340 360

large vs small scale anisotropy

averaged modulation over a given angular rangelow angular gradient

fine structurehigh angular gradientacceptance-corrected

4 hr = 60˚

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Paolo Desiati

2 hr = 30˚

4 hr = 60˚

IceCube

Milagro

20 TeV

1 TeV

10

Abbasi et al. (2011)

statistical significanceequatorial coordinates


Abdo et al. (2008)


Paolo Desiati

origin of small scale anisotropy ?astrophysics

Salvati & Sacco, arXiv:0802.2181Drury & Aharonian, Astropart. Phys. 29, 420 (2008)• CR from Geminga: ~90-200 pc, 340,000 yr ago

• magnetic connection & propagation in turbulent LIMF

• anisotropic MHD turbulence in the ISM

‣ particles streaming along magnetic field lines over ~100 pc (from a source) interact with O(1pc) ISM turbulence

‣ pitch angle scattering peaked near the direction of LIMF

Salvati, Astron. & Astrophys. arXiv:1001.4947

Malkov et al., ApJ 721, 750, 2010

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Paolo Desiati

origin of small scale anisotropy ?effect of turbulence

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‣ diffusion regime breaks down within mean free path

‣ interaction with turbulent interstellar magnetic field

‣ assuming an underlying dipole anisotropy, fractional localized regions form the effect of magnetic field turbulence

‣ the residual maps provide an image of magnetic field turbulence < 10’s pc

‣ cosmic ray energy spectra might also be affected by this propagation effects

Giacinti & Sigl, arXiv:1111.2536

20° smoothing

10 PeV

10 PeV

50 PeV


Paolo Desiati

diffusive propagation models ...

• … assume uniform diffusion coefficient across the Galaxy

• … do not account for energy-dependent interaction with ISM turbulence

• … do not account for magnetic field geometry

• … cannot explain non-dipolar anisotropy structures

• … break down within mean free path

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Paolo Desiati

from the Galaxy to our local interstellar medium

Milky Way

Benjamin (2008)

Frisch

Frisch

< 30,000 pc >

< 500 pc >

< 10-50 pc >

< 0.001 - 0.05 pc >

Local Bubble

Local Interstellar Cloud

Heliosphere

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Paolo Desiati

100 AU100 AU > 1,000 - 10,000 AU

~30

0-60

0 A

U

~1.4 TeV ~1.4 TeV ~14 - 140 TeV

~ 5

- 1

0 Te

V

3 µG

15

Rg ∼10−3

Z

�E

1 TeV

� �µG

B

�pcthe heliosphere

and the LIMF

Pogorelov & Zank (2004)

Vinterstellar flow ~ 26 km/s ≳ VAlfén


Paolo Desiati16

the heliospheremagnetic structure3D simulations of heliosphere Opher et al., arXiv:1103.2236 !

3D simulation of heliosphere/heliotailPogorelov et al., ApJ, 696, 1478, 2009

~150 AU ~1,000’s AU

~200 AU~0.1-1 AU


Paolo Desiati

• the wake downstream the interstellar flow develops turbulence from plasma velocity difference across the heliopause (similar to Kelvin-Helmholtz instability)

• charge-exchange processes decelerate the solar wind near the heliopause, producing an effective drag force that pushes the higher ISM density into the heliosheath. This generates Rayleigh-Taylor instability oscillations with amplitude 10’s AU over 100’s years - Liewer et al. (1996).

• charge-exchange processes in plasma-neutral fluid model produces alternate growing and damping of Alfvénic, fast and slow turbulence modes, with amplitude 10-100 AU and slowly propagating downstream along the heliopause - Shaikh & Zank (2010).

‣ The 10-100 AU turbulent ripples propagate outward the ISM and are damped by ion-neutral collisions in mfp ~ 300 AU - Spangler et al. (2011).

the heliosphereturbulence

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Paolo Desiati

scattering on heliosphericturbulence

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Paolo Desiati


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Paolo Desiati

• cosmic rays > 100 TeV do not feel the influence of the heliosphere

• cosmic rays < 100 TeV are influenced by the heliosphere from the downstream region

• resonant scattering of 1-10 TeV cosmic rays with 100’s AU turbulence ripples re-organizes the arrival direction distribution

• cosmic rays streaming along the LIMF experience the largest effect from the downstream region, and a minimal effect upstream

• perpendicular scattering is critical and determines the gradient region in cosmic ray arrival direction distribution

‣ evaluations and calculations to verify this scenario


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Paolo Desiati21

PD & Lazarian, arXiv:1111.3075

magneticequator

Funsten et al. (2009)Schwadron et al. (2009)Heerikhuisen et al. (2010)

LIMF direction compatible with• Ca II absorption & H I lines, Frisch (1996)• radio emission from inner heliosheath, Lallement et al. (2005), Opher et al. (2007)• polarization measurements, Frisch (2010)



Paolo Desiati22

PD & Lazarian, arXiv:1111.3075

magneticequator

Funsten et al. (2009)Schwadron et al. (2009)Heerikhuisen et al. (2010)

LIMF direction compatible with• Ca II absorption & H I lines, Frisch (1996)• radio emission from inner heliosheath, Lallement et al. (2005), Opher et al. (2007)• polarization measurements, Frisch (2010)


!


Paolo Desiati

spectral feature associated to anisotropyAbdo A.A. et al., Phys. Rev. Lett., 101, 221101 (2008)

harder spectrum in region A

2 hr = 30˚ 2 hr = 30˚

Milagro & ARGO-YBJ

harder than average spectrum from region A

γ < 2.7 at 4.6 σ levelEc = 3 - 25 TeV

similar to hardening of “diffuse” cosmic rays by Pamela, CREAM, ATIC-2

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Paolo Desiati

origin of spectral hardening ?

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‣ magnetic polarity reversals due to the 11-year solar cycles compressed by the solar wind in the magneto-tail

‣ turbulence makes reconnection fast and not affected by ohmic dissipation

‣ magnetic mirror @ single reconnection as site of acceleration (test particle)

Lazarian & PD, ApJ, 722, 188, 2010

!Sweet (1959) Parker (1957)

Lazarian & Vishniac, ApJ, 517, 700, 1999


Paolo Desiati

origin of spectral hardening ?

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‣ magnetic polarity reversals due to the 11-year solar cycles compressed by the solar wind in the magneto-tail

‣ turbulence makes reconnection fast and not affected by ohmic dissipation

‣ magnetic mirror @ single reconnection as site of acceleration (test particle)

Lazarian & PD, ApJ, 722, 188, 2010

~ 0.5 - 6 TeV


Paolo Desiati

Conclusions

• < 100 TeV cosmic ray anisotropy generated by interaction with the very local interstellar medium

• scattering with turbulence inside and in the outer heliospheric boundary to play an important role to explain large scale and small scale TeV cosmic ray anisotropy

• might explain change of cosmic ray anisotropy between 20 TeV and 400 TeV

• spectral hardening observed by Milagro & ARGO-YBJ from the downstream direction from re-acceleration of a fraction of cosmic rays in stochastic magnetic reconnection within the heliotail

• similar hardening observed by Pamela and CREAM could be related to the heliotail, although astrophysical explanations @ source and from propagation are possible

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