Cosmic Rays and Particle Acceleration

Nick Murphy

Harvard-Smithsonian Center for Astrophysics

Astronomy 253: Plasma Astrophysics

April 14–16, 2014

These lecture notes are based off of P. Carlson (2012), Zweibel (2013), Kulsrud, and slides by P. Blasi andothers. The derivations for first and second order Fermi acceleration follow High Energy Astrophysics by

Longair (3rd ed). This book contains several very nicely written chapters on relativistic particle dynamics,cosmic rays, and particle acceleration.

Outline

I History and importance of cosmic rays

I Detecting cosmic rays

I Properties of cosmic rays

I Second order Fermi acceleration

I First order Fermi acceleration (diffusive shock acceleration)

I Ultra-high energy cosmic rays

Introduction

I Cosmic rays (CRs) are highly energetic particles that mostlyoriginate outside of our solar system

I Cosmic rays are in rough energy equipartition with ourgalaxy’s magnetic field, the cosmic microwave background,and starlight

I The most fundamental questions of cosmic ray research are:I What are the source regions of cosmic rays?I How are cosmic rays accelerated?I How do cosmic rays propagate in the galaxy?

I Particle astrophysicists have made progress on all of thesequestions but much work remains

I Galactic and extragalactic cosmic rays are not the onlyenergetic particles that we care about: there are also solarenergetic particles (SEPs)!

History of cosmic ray research

I Research question in early 1900s: what is the source ofatmospheric ionization?

I Experiments by Coulomb in 1785 showed that a chargedmetallic sphere left alone in air gradually loses its charge

I Cosmic rays were first discovered by Victor Hess in balloonexperiments in 1912

I The amount of atmospheric ionization increased with heightI Origin of atmospheric ionization is extraterrestrial!

I Sea voyages by Clay (1927) and Millikan (1932) showed adependence of cosmic ray flux on latitude

I Modulation by Earth’s magnetic fieldI Effect more pronounced for lower energy cosmic raysI Cosmic rays are (mostly) charged particles!

I More CRs came from the west than the east (late 1920s to1930s)

I Most CRs are positively charged!

Cosmic ray research led to the discoveries of several newparticles

I Before particle accelerators, CRs were the best way to studyhigh energy particle physics!

I Cloud chamber observations of CRs by Anderson led to thediscovery of the positron (1932) and the muon (1936)

I Pions were discovered in CR observations in 1947

I The first CR antiproton was discovered in 1979

I These are secondaries resulting from collisions of primarycosmic rays with other nuclei

What effects do cosmic rays have on astrophysicalplasmas?

I Vertical support of ISM in galaxiesI Ionization and heating of ISM through collisions

I Allows weak ionization in molecular clouds, protoplanetarydisks, and the cold neutral medium

I Impacts chemistry in the ISM and molecular cloudsI Low energy CRs (. 10 MeV) are most responsible for this

I Probable driver of galactic windsI Amplification of magnetic fields

I Upstream and downstream of shocksI Cosmic ray driven dynamo

I Modification of astrophysical plasma processesI Shocks, dynamos, jets, reconnection

Why should humans care about cosmic rays?

I Contribution to atmospheric chemistryI Production of 14C in the atmosphere through the neutron

capture reactionn + 14N→ 14C + p (1)

I Historical spikes of 14C production (e.g., 774–775 CE) possiblydue to cosmic ray events

I Source of background radiationI Airline pilots have a higher risk of cataracts due to CRs

I CRs and SEPs are a health hazard for interplanetary spacetravel (in particular, outside of Earth’s magnetosphere)

I Impact on Earth and space based electronicsI Cosmic rays cause bit flips (e.g., error on Voyager 2)I Increases with altitude in atmosphereI Error correction schemes can mitigate these effects

Why should particle physicists care about cosmic rays?

I The maximum energy of cosmic rays is of order 1020 eVI Orders of magnitude greater energy than accessible by the

Large Hadron Collider (LHC) at CERN!I Possible to investigate particle physics at otherwise inaccessible

energies

I Higher energy cosmic rays induce particle showers inatmosphere

I Difficulties:I Lack of experimental controlI Lack of cross-checks with experiment at these energiesI Rarity of most energetic events

Primary cosmic rays collide with molecules in theatmosphere and yield secondary particles in an air shower

I Common products: pions, electrons, positrons, neutrons,kaons, muons, and neutrinos

I Particles produce Cherenkov radiation that is detectable

Particle profile: pions

I Pions are mesons: they consist of a quark and antiquark

π+ : ud π0 : uu or dd π− : du

I Charged pions have a mass of 139.6 MeV/c2. The mostcommon decay routes are

π+ → µ+ + νµ (2)

π− → µ− + νµ (3)

with a characteristic decay time of 26 ns

I Neutral pions have a mass of 135.0 MeV/c2 and usually decayinto two photons with a characteristic time of 8.4× 10−17 s:

π0 → 2γ (4)

I Gamma rays from pion decay are detectable by Fermi

Particle profile: muons

I Muons (µ−, µ+) are leptons with a mass of 105.7 MeV/c2

(about 200 times the mass of an electron)

I Most naturally occuring muons are cosmic ray secondariesresulting from pion decay

I Muons have a mean lifetime of 2.2 µsI That’s really long!I Time dilation allows muons to propagate longer distances in

our frame before decaying

I The most common decay routes are

µ− → e− + νe + νµ (5)

µ+ → e+ + νe + νµ (6)

How are cosmic rays observed?

I Synchrotron radiation

I Gamma rays (pion decay bump)I Cloud chambers and bubble chambers

I Important historically; mostly obsolete

I Magnetic spectrometersI High precision measurements of bending of cosmic ray paths

by an applied magnetic fieldI Useful at relatively low energies

I Cherenkov radiationI Particles traveling faster than the speed of light in a medium

Alpha Magnetic Spectrometer

I A magnetic spectrometer1 on the International Space Station

I Above atmosphere so it measures primary cosmic raypopulation

I Measuring electron-to-positron ratioI Provides limits on annihilation of dark matter particles

I Required a special act of Congress to get it launched!

1The first rule of tautology club is the first rule of tautology club.

Pierre Auger Observatory

I Observes air showers of secondary cosmic ray particles

I Rate of ultra-high energy cosmic rays is ∼1 km−2 century−1

I Effective area is ∼3000 km2

I Fluorescence detectors to observe air showers in atmosphere

I Cherenkov detectors to detect secondary particles on ground

I Combination of both methods reduces systematic errors thatresulted during prior work

I A key result: anisotropy of UHECRs (more later)

What are the properties of cosmic rays?

I Energy spectrum

I Abundances

I Confinement/lifetimes

I Isotropy/anisotropy

Cosmic ray research is a great way to get practice withrarely used SI prefixes!

Energy spectrum of cosmic rays

Energy spectrum of cosmic rays

I Knee around 3× 1015 eV

I Ankle around a few × 1018 eV

Energy spectrum of cosmic rays

I Flux of lowest energy cosmic rays is uncertain due toheliospheric modulation

I Solar energetic particles present at lower energies (∼keV to∼100 MeV)

I From ∼109 to ∼1015 eV the power law index is ≈2.7I This part of the spectrum contains most of the cosmic ray

energy density

I Above the knee the index steepens to ≈3.0

I The spectrum starts to flatten again around the ankle at afew × 1018 eV

Why do the knee and ankle exist in the cosmic ray powerspectrum?

I Changes in the spectral index may indicate a transition in theacceleration or confinement mechanism and/or composition

I The knee is postulated to be an upper limit associated withthe acceleration mechanism for galactic cosmic rays

I Can diffusive shock acceleration in supernova remnants reachthis energy?

I Does the knee have anything to do with particle transport orescape?

I Cosmic rays above the ankle are extragalactic in originI Gyroradius is comparable to size of galaxy

I Neither the knee nor the ankle is well-understood

I Question: how could we determine if cosmic rays between theknee and ankle are galactic or extragalactic?

Abundances of cosmic rays

Abundances of cosmic rays

I Relative abundances are more or less representative ofinterstellar plasmas

I H and He are somewhat underabundant

I ∼89% are protons

I ∼10% are α particles

I Overabundance of 3He, Li, Be, and B (by 5–7 orders ofmagnitude at ∼1 GeV)

I 1–2% are electronsI Steeper slope because of radiative lossesI Very low density of cosmic rays → maintaining quasineutrality

is not a problem

I Very small fraction of CRs are positrons or antiprotons

I No antihelium nuclei yet discovered in cosmic rays

Are Coulomb collisions significant?

I Recall: Coulomb collisions occur when a particle interactswith another particle via electric fields

I The Coulomb cross section of a 1 GeV particle is 10−30 cm2:tiny!

I For 1 GeV cosmic rays propagating in the ISM (n ∼ 1 cm−3),the mean Coulomb collision rate is

nσv ∼ 10−19.5 s−1 (7)

This corresponds to a 1% chance of a Coulomb collisionduring a Hubble time!

I Coulomb collisions can therefore be neglected completely

Spallation reactions provide limits on cosmic ray lifetimes

I Species such as 3He, Li, Be, and B are very overabundant ingalactic cosmic rays

I Spallation occurs when C, N, O, Fe nuclei impact interstellarhydrogen

I The larger nuclei are broken up into smaller nucleiI The abundance enhancements of the lighter elements show

that spallation reactions are somewhat importantI Must have passed through & 3 g cm−2 of the ISM

I Heavier elements such as Fe are only mildly depletedI Cannot have passed through more than ∼5 g cm−2 of material

I This corresponds to a cosmic ray lifetime of 3 Myr in the disk(or significantly longer in the halo)

I Radioactive isotopes provide an additional diagnosticI Example: 10Be has a half-life of 1.5 Myr

Cosmic rays are effectively confined if the Larmor radius ismuch smaller than the system size

I Relativistic momentum: p = γmv

I The relativistic Larmor radius is

rL =p⊥|q|B

=γmv

ZeB(8)

This can be written for a proton as

rL =pc/eV

300(B/G)cm (9)

How effectively are cosmic rays confined in differentsystems?

I Calculate the particle energy associated with a gyroradiusequal to the characteristic length scale

I Heliosphere: L ∼ 100 AU, B ∼ 10 µG ⇒ EL ∼ 4× 1012 eVI Low energy cosmic rays are modulated by heliospheric

magnetic fields

I ISM: L ∼ 100 pc, B ∼ 5 µG ⇒ EL ∼ 4× 1017 eVI The highest energy cosmic rays are barely deflected by the

magnetic field of the Milky Way

Isotropy of cosmic rays below the knee

I Below ∼1014 eV, cosmic rays are highly isotropic

I Cosmic rays propagate along field lines in galaxy

I Propagation and confinement make it difficult to determinesource regions

I Scattering and diffusion play an important role

I Weak anisotropy (above) is not well understood

What properties of cosmic rays must an accelerationmechanism explain?

I A power law energy spectrum for particles of all types:

dN(E ) ∝ E−xdE (10)

The exponent x is usually in the range of ∼2.2–3.0.I The acceleration of cosmic rays to maximum observed

energiesI For galactic cosmic rays, energies up to the knee: ∼1015 eVI For extragalactic cosmic rays, energies beyond the ankle:∼1020 eV

I Elemental abundances of cosmic rays similar tointerstellar/circumstellar abundances

General principles of acceleration

I The equation of motion for a charged particle is

d

dt(γmv) = q

(E +

v × B

c

)(11)

where γ ≡ 1√1−v2/c2

is the Lorentz factor and q, m, and v are

the charge, mass, and velocity of the particle

I Magnetic fields themselves do no work and cannot be directlyresponsible for acceleration, but changing magnetic fields leadto an inductive electric field

Direct acceleration by a static electric field

I Direct acceleration requires a nonzero average electric field:

〈E〉 6= 0 (12)

I Electric fields that develop due to charge separation willquickly short themselves out by motion of free charges

I Thus, we’d normally expect that direct acceleration will notsuffice to describe observations of CRs

I Important exception: magnetic reconnectionI Development of strong localized E

Second order Fermi acceleration2

I Prior to 1949, suggested sources of cosmic rays includedmagnetically variable stars and supernovae

I Details of proposed mechanisms were lackingI Difficulty in explaining the observed cosmic ray spectrum

I Fermi (1949) proposed a model for acceleration in whichparticles can statistically gain energy through collisions withinterstellar clouds

I Original model led to second order acceleration: ∆EE ∝

(Vc

)2

I By assuming a characteristic escape time, τesc , this results in apower law distribution of particle energies

2Following Longair’s textbook

Acceleration occurs by particles having random collisionswith interstellar clouds with characteristic velocity V

I Reflection of particles due to magnetic mirror effects

Geometry of collisions

I Energy is gained during a head-on collision (a) and lost duringa trailing collision (b)

I Parameters:I V : characteristic cloud velocity in observer’s frameI v : velocity of particleI θ: angle of incidenceI m: particle mass (� M ≡ cloud mass)

Key point: head-on collisions are more likely than trailingcollisions

I The probability of a head-on collision is proportional to

v + V cos θ (13)

I The probability of a trailing collision is proportional to

v − V cos θ (14)

I Same principle as driving on a highway: there will be morecars passing you by from oncoming traffic than traffic going inthe same direction

I Further generalizable to 3D bumper cars piloted by spacewombats

It is necessary to do a fully relativistic analysis

I The particle energy E in the cloud’s (center of momentum)reference frame is

E ′ = γV (E + Vp cos θ) (15)

where p is the particle’s momentum, θ is the angle ofincidence, and the Lorentz factor is

γV =

(1− V 2

c2

)−1/2

(16)

Reversing the x component of momentum in cloud frame

I The x component of momentum in this frame is

p′x = p′ cos θ′ = γv

(p cos θ +

V Ec2

)(17)

I During the collision:I The particle’s energy is conservedI The x component of the particle’s momentum is reversed

p′x → −p′x (18)

I Then transform back into the observer’s frame:

E ′′ = γV(E ′ + Vp′x

)(19)

Finding the change in energy for a collision

I Combinging Eqs. 15, 17, & 19 yields an expression for theparticle energy in the observer’s frame:

E ′′ = γ2E

[1 +

2Vv cos θ

c2+

(V

c

)2]

(20)

I Then expand to second order in V /c and solve for ∆E

E ′′ − E ≡ ∆E = E

[2Vv cos θ

c2+ 2

(V

c

)2]

(21)

We still need to average over the angle of incidence, θ

I Assume that particle is randomly scattered in pitch angleI Distribution of angles is random

I Simplify by assuming that the particles are relativistic: v ≈ cI Use that the probability of an angle between θ and θ + dθ is

proportional to sin θ dθI Analogous to the reason why we see more galaxies edge-on

than face-on

I Average over the range 0 to π

I The average of the θ-dependent term in Eq. 21 becomes⟨2V cos θ

c

⟩=

(2V

c

) ∫ 1−1 X

[1 + V

c X]dX∫ 1

−1

[1 + V

c X]dX

=2

3

(V

c

)2

(22)

where we use the substitution X ≡ cos θ

The average increase of energy per collision is only secondorder in V /c

I The full average over theta finally yields and expression for theaverage increase in energy during a collision⟨

∆EE

⟩=

8

c

(V

c

)2

(23)

I The same fractional increase occurs during every collision

I Define L as mean free path and φ as pitch angle

I The time between collisions ∼ L/c cosφ, which can beaveraged to 2L/c

I The average rate of energy increase is

dEdt

=4

3

(V 2

cLE)

= αE (24)

What is the spectrum of the accelerated particles?

I The diffusion-loss equation provides a convenient way ofcalculating the spectrum

dN

dt= D∇2N +

∂

∂E[b(E)N(E)]− N

τesc+ Q(E) (25)

I N(E) is the number of particles between E and E + ∆EI D∇2n is the diffusive term (which we ignore)I b(E) = −dE/dt is the energy loss termI Q(E) represents sources (which we ignore)I τesc is the escape time

Deriving the energy spectrum

I If we assume a steady state and use our expression for a gainin energy, b(E) = −αE , then this equation becomes

− d

dE[αEN(E)]− N(E)

τesc= 0 (26)

I Differentiating then yields

dN(E)

dE= −

(1 +

1

ατesc

)N(E)

E(27)

I The solution is thenN(E) ∝ E−x (28)

where x ≡ 1 + (ατesc)−1

I This is a power law! But α and τesc are model-dependent.

What are the problems with second order Fermiacceleration?

I The random velocities of clouds are relatively small:V /c . 10−4

I For a CR mean free path of ∼0.1 pc, collisions would likelyoccur only a few times per year

I Very little chance of a particle gaining significant energy!

I This theory does not include energy losses that compete withenergy gains (e.g., ionization)

I The theory does not predict the power law exponentI Why are the spectra of energetic particles so similar?

I Biggest problem: energy gains are second order: O(V 2/c2)I Some collisions result in energy losses!

Setting the stage for first order Fermi acceleration

I What if we could find a similar mechanism in which everyinteraction resulted in an increase in energy?

I This would be a first order process! O(V /c)

I The occurs during diffusive shock acceleration if particles canbounce back and forth in the upstream and downstreamregions and always approach plasma moving toward them

I Not quite the same idea as improving gas mileage by makingall roads downhill

I An energetic particle propagating through the shock frontwould barely notice it if the particle’s gyroradius is muchlarger than the shock thickness

I The velocity becomes isotropized due to streaminginstabilities and turbulent motions

I Define U as the shock velocity

In the pre-shock and post-shock reference frames

I (a) Observer’s frame, (b) reference frame of shock,(c) upstream frame, (d) downstream frame

I When crossing the shock from either side, the particle seesplasma moving toward it at a velocity of V ≡ 3

4U

Let’s look at Fermi acceleration more generally

I Define E = βE0 as the average energy of the particle after acollision

I Define P as the probability that the particle remains in theacceleration region after a collision

I After k collisions there are N = N0Pk particles with energies

E = E0βk

Relating the spectrum to energy gain and loss probability

I Eliminating k yields

ln(N/N0)

ln(E/E0)=

lnP

lnβ(29)

⇒ N

N0=

(EE0

)lnP/ lnβ

(30)

I The power spectrum is then

N(E)dE ∝ E−1+(lnP/ lnβ) dE (31)

I From second order Fermi acceleration, this means that

lnP

lnβ≡ −(ατesc)−1 (32)

Let’s evaluate the fractional increase in energy for oneacceleration cycle

I A particle with energy E on the upstream side sees plasmaapproaching it at a velocity V from the downstream side

I The particle’s energy in the downstream frame is

E ′ = γV (E + pxV ) (33)

I Assume shock is non-relativistic so γV = 1

I The particles are relativistic so px = Ec cos θ and therefore

∆E = pV cos θ (34)

∆EE

=V

ccos θ (35)

where we use that E = pc

What is the probability of an angle θ?

I The probability that a particle will cross the shock with anangle of incidence between θ and θ + dθ is proportional tosin θ dθ

I The rate at which they approach the shock front isproportional to vx = c cos θ

I For θ ∈ [0, π/2], the probability function for θ is

p(θ) = 2 sin θ cos θ dθ (36)

I Eq. 35 then averages to⟨∆EE

⟩=

V

c

∫ π/2

02 cos2 θ sin θ dθ

=2

3

V

c(37)

which gives the energy increase for one shock crossing

Because plasma appears to be moving towards particlescrossing the shock front at equal velocities from each side,the energy gain is symmetric

I The fractional increase in energy after a round trip is given by⟨∆EE

⟩=

4

3

V

c(38)

I This means that the energy change after a round trip is givenby

β ≡ EE0

= 1 +4

3

V

c(39)

I First order Fermi acceleration is first order in Vc !

How do we estimate the escape probability per cycle?

I The rate of particles crossing the shock from either directionis 1

4nc , where n is the number density of energetic particlesand we average over velocities/angles

I Upstream particles are swept into the shock, so few losses

I Particles removed from downstream region at a rate ofnV = 1

4nU

I The fraction of particles lost is then 14nU/

14nU = U/c so that

the fraction of particles remaining after one cycle is

P = 1− U

c(40)

Since U � c, few particles are lost per cycle

What is the power law exponent?

I We then have

lnP = ln

(1− U

c

)= −U

c(41)

lnβ = ln

(1 +

4

3

V

c

)' 4

3

V

c=

U

c(42)

so thatlnP

lnβ= −1 (43)

I The energy spectrum is then

N(E) dE ∝ E−2 dE (44)

I This derivation predicts a value of the power law exponentwhich is not too different from the observed value of ∼2.7 forgalactic cosmic rays

First order Fermi Acceleration

I This model provides an efficient method for acceleratingparticles in supernova remnant shock waves

I Requires that velocity vectors are randomized upstream anddownstream of the shock

I Scattering of CRs by Alfven waves, etc.

I Requires prior injection of superthermal seed particles

I Energy losses are ignored

I This model neglects detailed physics of the magnetic field

I Key paper: Bell (1978). See also Longair book.

What is the maximum attainable energy by thismechanism?

I The magnetic field must be able to confine energetic particles

I For interstellar magnetic field strengths ∼ 1µG, the protonLarmor radius is ∼ 1 pc

I Supernova remnants could then not confine particles withenergies &1014 eV

I But what if there are mechanisms to amplify the magneticfield?

Magnetic perturbations can be amplified by super-Alfvenicstreaming of cosmic rays

Magnetic field amplification by cosmic rays

I Nature of instabilities (e.g., Zweibel 1978; Bell 1978, 2004)I Upstream plasma responds to cosmic ray current with return

currents to compensate for positive chargeI The induced perturbations grow

I Key questions about these instabilities:I What is the nonlinear evolution and saturation of these

instabilities?I How do cosmic rays diffuse?

I Can explain acceleration to 1017–1018 eV in supernovaremnant shock waves, but not to 1020 eV

I If we had more time, we could have an entire class on thesetopics!

Observation of the pion decay bump in a supernovaremnant by Fermi (Ackermann et al. 2013)

I 67.5 MeV in rest frame of π0

I Detection difficult because energetic electrons also emitgamma rays through bremsstrahlung and inverse Comptonscattering

Open questions on galactic cosmic rays

I Why does the energy spectrum change its slope at ∼1015 eV?

I How are cosmic rays accelerated to ∼1017–1018 eV?

I Are cosmic rays between the knee and ankle: galactic,extragalactic, or a combination?

I What is the nature of cosmic ray diffusion and transport ingalactic magnetic fields?

I What is the nature of instabilities that connect the magneticfield and cosmic ray populations?

I Especially the nonlinear evolution

Ultra-high energy cosmic rays

I Major problem: how are particles accelerated to 1020 eV?I The energy of a fastball in a single particle!

I Orders of magnitude more energy than any particleaccelerator humans have constructed!

I Venturing into regions of particle physics parameter space thatcannot be cross-checked in the laboratory

I Cannot be confined in the Milky WayI Extragalactic source

I What is the nature of the primary particles?

I State-of-the-art observations are from the Pierre AugerObservatory

Possible sources of UHECRs (‘Hillas plot’ from Kotera &Olinto 2011)

I UHECRs of these types can be confined above these lines

Possible sources of UHECRs

I Mechanisms includeI Fermi acceleration at shock waves (mildly relativistic shocks?)I Unipolar inductors (neutron stars?)

I Possible source regions includeI Gravitational accretion shocks (around clusters of galaxies?)I Active galactic nuclei (jets? central region?)I Gamma ray bursts (highly magnetized shocks?)I Neutron stars (magnetars?)

I Much work needs to be done both theoretically andobservationally

Intergalactic propagation of cosmic rays

I At energies above the knee, pair production becomes possiblethrough interaction with cosmic microwave background(CMB) photons

p + γCMB → p + e+ + e− (45)

I Fraction of energy lost in each interaction is small so effectiveinteraction distance remains large: ∼ 600 Mpc

I At energies above ∼1020, pion production becomes possible

p + γCMB → p + π0 (46)

p + γCMB → n + π+ (47)

I More energy is lost so the interaction distance is muchshorter: 50–100 Mpc (GZK limit)

I UHECRs are consequently thought to originate closer thanthis distance

I Are there UHECRs above this limit from further away?

Anisotropy of ultra-high energy cosmic rays

I The Pierre Auger Observatory has measured anisotropy ofUHECRs (above 5× 1019 eV)

I Circles indicate UHECRs; red crosses indicate nearby AGN

I Arrival directions correlated with AGN within ∼75 MpcI Identification is tentative

I Could this actually be the last scattering surface instead of thesource?

Solar Energetic Particles (SEPs)

I Before and after a solar eruption that led to an SEP event

Solar Energetic Particles (SEPs)

I Two sources that result from solar flares, coronal massejections, and other solar eruptions

I Reconnection (impulsive events)I Shocks (gradual events)

I Energies up to ∼100 MeV or ∼1 GeVI SEPs propagate mostly along field lines

I We’ll observe an SEP event if there are magnetic field linesattaching the source region to us

I Occasional events are observed at spacecraft 180◦ apart

I One of the major components of space weather

I Observations of solar flares suggest that ∼all of the electronsin the source region are accelerated

Acceleration due to contracting islands during reconnection

I ‘Magnetic islands’ often form during reconnectionI ‘Flux ropes’ in 3D

I Drake et al. have suggested that contraction of these islandsleads to Fermi acceleration of confined particles

I Applied to solar flares, boundary of heliosphere, etc.

I Alternative to direct acceleration by the electric field

I 3D structure is vital to how this process works

The Sun emits gamma rays!

I Two source regions:I Compact source: cosmic ray cascades in photosphereI Extended source: inverse Compton scattering in corona during

which cosmic rays give a kick to photons

Summary

I Cosmic rays are important astrophysically, terrestrially, and inparticle physics

I We can actually detect stuff besides photons!

I The most important observed/inferred properties of cosmicrays include their energy spectrum, abundances,confinement/lifetimes, and isotropy

I Diffusive shock acceleration in supernova remnants is themost promising mechanism for cosmic ray acceleration toenergies up to the knee (. 1015 eV)

I Cosmic-ray driven instabilities amplify the magnetic fieldsurrounding shocks in supernova remnants

I Improves confinement; allows acceleration to higher energies

I Summary of UHECRs: we don’t know their composition,where they come from, or how they were accelerated