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Berlin, 2009 29 Sept Tom Gaisser 1 Cosmic Rays 1. Introduction Lecture 1: Introduction to cosmic rays Lecture 2: Atmospheric μ and ν and ν telescopes Lecture 3: Giant air shower experiments Exercise: Expectations for γ-ray & ν-astronomy

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Cosmic Rays 1. Introduction

Lecture 1: Introduction to cosmic raysLecture 2: Atmospheric µ and ν and ν telescopesLecture 3: Giant air shower experimentsExercise: Expectations for γ-ray & ν-astronomy

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Highlights of history of cosmic rays

• 1912: Victor Hess, Nobel prize 1936– Ascended to 5300 m in an open balloon– Showed ionization of the air increased

• ~1920: Millikan (Caltech)– coined “cosmic rays” – thought they were photons– Compton (at Chicago) argued they were positive particles

• 1930’s: latitude surveys, Geiger counters– E-W effect proves primary cosmic rays are +charged– Atmospheric secondaries are muons, photons and e±

• 1940’s: (post-war) photographic emulsions– 1947 Powell et al., discovery of the pion– 1948 discovery of nuclei in primary cosmic radiation– 1949 Fermi (Chicago) paper on cosmic-ray acceleration

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Air shower history

• 1937: Pierre Auger discovery– Observed earlier by Rossi

• 1940’s: Rossi (M.I.T.) – Air shower studies with arrays of scintillators

• 1960’s: discovery of the “knee” – Bernard Peters points out need for a new population– 1962: John Linsley (Volcano Ranch)

• "Evidence for a Primary Cosmic-Ray Particle with Energy 1020 eV". Physical Review Letters 10: 146, 1963

• AGASA, Fly’s Eye, HiRes, Auger, TA …

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Tibet hybrid air shower array4300 m ~600 g/cm2

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125 m

Photo by James Roth17 December 2007

IceTop

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Cosmic ray showers from above

Neutrinos from all directions

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γ-ray telescopes

H.E.S.S.

VERITAS

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CR Iron spectrum from Direct Cherenkov

HESS collaboration, Phys Rev D 27, 2007

Method proposed by Kieda et al, 2001

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Rigidity-dependence

• Acceleration, propagation– depend on B: rgyro = R/B– Rigidity, R = E/Ze– Ec(Z) ~ Z Rc

• rSNR ~ parsec – Emax ~ Z * 1015 eV– 1 < Z < 30 (p to Fe)

• Slope change should occur within factor of 30 in energy

• With characteristic pattern of increasing A

• Problem: continuation of smooth spectrum to EeV

• More on this later…

30

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B. Peters: if Emax depends on Bthen p disappear first, then He, C, O, etc.

B. Peters, Nuovo Cimento 22 (1961) 800

Peters cyclePeters cycle: systematic increase of < A > : systematic increase of < A > approaching approaching EEmaxmax

<A> should begin to decrease again<A> should begin to decrease againfor E > 30 x for E > 30 x EEkneeknee

gyro-radius = Pc / ZeB ≡ R (rigidity) / B Etotal (knee) ~ Z x R(knee)

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Spectrometers(∆A = 1 resolution,good E resolution)

Calorimeters(less good resolution)

Primary spectrum

Airshowers

Air-shower arrayson the ground toovercome low flux.Don’t see primariesdirectly.

Current questions

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Scientific program

Solar and HeliosphericSH.1: SOLAR EMISSIONS

SH.2: ACCELERATION AND PROPAGATION IN THE HELIOSPHERE SH.3: GALACTIC COSMIC RAYS IN THE HELIOSPHERE

Origin and GalacticOG.1: LOW ENERGY COSMIC RAYS

OG.2: X-RAY, GAMMA-RAY AND NEUTRINO ASTRONOMY AND ASTROPHYSICS

High EnergyHE.1: HIGH ENERGY COSMIC RAYS – EAS

HE.2: PARTICLE PHYSICS, ASTRO-PARTICLE PHYSICS AND COSMOLOGY

31st Int. Cosmic Ray Conf.July 7-15, 2009, Łόdź, Poland

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Solar flare shock acceleration

Coronal mass ejectionCoronal mass ejection09 Mar 200009 Mar 2000

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Movie: LASCO on SOHO

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1998 April 11 - 20: An interesting period: Between April 10 and April 13, a comet enters the field of view from the left (East) and passes around the Sun. On April 10-11, a smaller sun-grazing comet approaches the Sun from the south, just to the right of the occulter pylon. The period culminates in a fast CME associated with a high energy particle storm at SOHO

Movie from LASCO instrument on SOHO

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13 Dec 2006 solar flare: GLE in IceTop

With 32 tanks in 2006

Ap. J. (Letters) 689 (2008) L65-L68

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Solar flare

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Lessons from the heliosphere

• ACE energetic particle fluences:• Smooth spectrum

– composed of several distinct components:

• Most shock accelerated• Many events with different shapes

contribute at low energy (< 1 MeV)

• Few events produce ~10 MeV

– Knee ~ Emax of a few events– Ankle at transition from heliospheric

to galactic cosmic rays

R.A. Mewaldt et al., A.I.P. Conf. Proc. 598 (2001) 165

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Heliospheric cosmic rays

• ACE--Integrated fluences:– Many events contribute to

low-energy heliospheric cosmic rays;

– fewer as energy increases.– Highest energy (75 MeV/nuc)

is dominated by low-energy galactic cosmic rays, and this component is again smooth

R.A. Mewaldt et al., A.I.P. Conf. Proc. 598 (2001) 165

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Measurements from TRACER, BESS and others

Galactic cosmic rays

Figure by P. Boyle & D. Müller in RPP

dN / dE ~ Eγγγγ ~ 2.7

-γγγγ

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BESS spectrometer

AΩ ~ 0.085 m2 sr

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A fundamental result

• Excess of Li, Be, B from fragmentation of C, O

• Spallation σ plus ρISM give dwell time of nuclei– Find τ ~ 3 x 106 yrs– cτ ~ Mpc >> size of

galactic disk (kpc)– Suggests diffusion in

turbulent ISM plasma– Predictions for γ-rays,

positrons and antiprotons follow

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Energy-dependence of secondary/primary cosmic-ray

nuclei • B/C ~ E-0.6

• Observed spectrum:– φ(E) = dN/dE ~ K E-2.7

• Interpretation:– Propagation depends on E

– τ(E) ~ E-0.6

– φ(E) ~ Q(E) x τ(E) x (c/4π)

• Implication:– Source spectrum:

Q(E) ~ E-2.1Garcia-Munoz, Simpson, Guzik, Wefel& Margolis, Ap. J. 64 (1987) 269

B / C

sub-Fe / Fe

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Secondary / primary ratios from CREAM

arXiv: 0808.1718δδδδ = 0.6

δδδδ = 0.33

5% primary N / O

15% primary N / O

• B is entirely secondary• Source is production in ISM• Spectral index at production

is same as observed(propagated) primary nuclei

• Therefore observed spectralof B is 2.7 + δ = 2.1 + 2δ ~3.3

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Energetics of cosmic rays• Total local energy

density: – (4π/c) ∫ Eφ(E) dE

~ 10-12 erg/cm3 ~ B2 / 8π• Power needed:

(4π/c) ∫ Eφ(E) / τesc(E) dEgalactic τesc ~ 107 E-0.6 yrsPower ~ 10-26 erg/cm3s

• Supernova power:1051 erg per SN~3 SN per century in disk~ 10-25 erg/cm3s

• SN model of galactic CRPower spectrum from shock

acceleration, propagation

Spectral Energy Distribution (linear plot shows most E < 100 GeV)

(4π/c) Eφ(E) = local differential CR energy density

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Cosmic rays in the Galaxy

• Supernova explosions energize the ISM

– ~1% Kinetic energy;

neutrinos ~ 99%

– >10% of kinetic energy CR acceleration

– Energy density in CR ~ B2/8π

– SN & CR activity drives Galactic wind into halo (Parker)

– CR diffuse in larger volume

– Eventually escape Galaxy

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Diffuse galactic secondaries

p + gas π0, π+/−, antiprotons

• π0 γ γ [π+/− ν, µ e+/−]• γ-spectrum depends on where

production occurs

BESS antiprotons, 1997, ’99, ’00.•Fully consistent with secondary production by collisions in ISMfollowed by solar modulationvarying with solar cycle

Phys.Rev.Lett. 88 (2002) 051101

Before Fermi & PAMELA

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Kinematics of π0 decay

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p p π0 + anything

γ γp p p + anything

π0 γ γp per proton interaction

Constructing secondary spectrumfrom primary spectrum

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Are all photons and p produced in ISM?

Secondary production: V. S. Ptuskin et al, ApJ 642 (2006) 902

Egret diffuse γγγγ .)PAMELA antiprotons (prelim.)

A fraction may be produced in the sources

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Diffiuse gamma-rays from FERMI

Luca Baldini (INFN) SOCoR, Trondheim June 18, 2009 –FERMI - preliminary

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Antiproton to proton ratio

PRL 102, 051101 (2009)Seconday Production Models

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Acceleration

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Next step is to average over cos θ2 and cos θ1

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Distribution of exit angles (cos θ2 )

cos θ2 averages to 0 so

Distribution of entrance angles (cos θ1 ):

cos θ1 averages to -V / 3c so

cos θ2 averages to 2/3 so that

then

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Supernova progenitor

SN ejecta Shocked ISM

Supernova blast wave acceleration

Unshocked ISM

SNR expands into ISMwith velocity V~ 104 km/s.Drives forward shock at 4/3 V

Forward shock

u1 ~ 4/3 V

u1 ~ 4/3 V

Particle with E1

E2 = ξ ξ ξ ξ E1

Contact discontinuity, V

TSN ~ 1000 yrs before slowdownEmax ~ Z x 100 TeV

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Problems of simplest SNR shock model

• Expected shape of spectrum:– Differential index α ~ 2.1 for

diffusive shock acceleration• αobserved ~ 2.7; αsource ~2.1;

∆α ~ 0.6 τesc(E) ~ E-0.6

• c τesc Tdisk ~100 TeV• Isotropy problem

• Emax ~ βshock Ze x B x Rshock

– Emax ~ Z x 100 TeV with exponential cutoff of each component

– But spectrum continues to higher energy:

• Emax problem

• Expect p + gas γ (TeV) for certain SNR– Need nearby target as shown

in picture from Nature (F. Aharonian, April 02)

– Some likely candidates (e.g. HESS J1745-290) but still no certain example

– Problem of elusive π0 γ-rays

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Solutions to problems ?

• Isotropy problem – γsource ~ 2.3 and τesc ~ E-0.33 ??– τesc ~ E-0.33 is theoretically preferred, but γ = 2.3?

• Evidence for acceleration of protons at SNR– Look for TeV γ-rays mapping gas clouds at SNR– Possible example SNR IC443 (VERITAS,

arXiv:0810.0799)

• Acceleration to higher energy– Magnetic field amplification in non-linear shock

acceleration

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Non-linear shock acceleration

• Discussion of acceleration so far– “Test particle” approximation – Neglects effects of particles being accelerated on the

accelerator• Non-linear theory

– Accounts for magneto-hydrodynamic turbulence generated by cosmic-rays upstream of the shock.

– Cosmic-rays scatter on MHD turbulence with wavelength matched to particle’s gyro-radius

• Two important effects– Spectrum is distorted– Magnetic field in acceleration region is amplified

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Non-linear shock acceleration - 2• Cosmic-ray pressure in upstream region generates precursor • Higher energy particles get further upstream• Experience larger discontinuity r = u1(x) / u2

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Non-linear shock acceleration - 3

• Spectrum at shock concave, • γ < 1 at high energy,• energy concentrated at Emax

Plot from P. Blasi

Note: CR acceleration theorycalculates distribution in p-space:f(p) ≡ dncr / dp3 ~ (1/p2) dn/dE

So here q = γ + 3

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Non-linear shock acceleration - 4

• Cosmic-ray pressure also amplifies magnetic field (Bell, MNRAS 353, 550, 2004)

• Emax increases to 1016 or 1017 eV in early,free-expansion phase of SNR expansion

• applies to small fraction of accelerated particles

ALSO (Ptuskin & Zirakashvili, A&A 429, 755, 2005 )• In later phases of SNR expansion:

- upstream scattering becomes inefficient as expansion slows down

- Emax Emax (t) decreases with time- Accelerated particles with E > Emax(t)escape upstream

- observed spectrum is integral of SNR history

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ChandraCassiopeia A

ChandraSN 1006Filamentary structure of X-ray emission

of young SNRs:Evidence for amplification, B ~ 100 µµµµG

Berezhko&Volk, OG111

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RX J1713.7-3946Berezhko & Völk, arXiv:0707.4647

Contributions from electrons (IC, NB) suppressed by ~100 µµµµG fields

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Examples of power-law distributions(M.E.J. Newman, cond-mat/0412004)

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More examples from M.E.J. Newman, cond-mat/0412004

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Casualties per attack in Iraq(Neil F. Johnson, et al., from APS News, 8 Nov 2006)

Differential α ~ 2.5

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Frank Capra, 1957


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