COMPOSITION AND PROPAGATION OF GALACTIC COSMIC RAYSBELOW THE KNEE

DIETRICH MÜLLERUNIVERSITY OF CHICAGO

PAMELA PHYSICS WORKSHOPROME, MAY 12, 2009

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OUTLINE1. Introduction: Comments on History2. Cosmic-Ray Sources: Observational Constraints3. The Consensus Model4. The Experimental Challenge5. The TRACER Approach6. A Propagation Model7. What about Protons and Electrons?8. Conclusions

1. At which Energies to look?

2. Experimental Challenges3. The TRACER Approach4. A self-consistent Model5. What is still needed?6. Conclusions

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HISTORY

Baade and Zwicky 1934Supernova Explosions

But how does it work??? No magnetic Fields?

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HISTORY

Baade and Zwicky 1934Supernova Explosions Fermi 1949

Distributed Acceleration But how does it work??? No magnetic Fields?

Efficiency???Does it work for heavy Elements?

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HISTORY

Baade and Zwicky 1934Supernova Explosions Fermi 1949

Distributed Acceleration But how does it work??? No magnetic Fields?

Efficiency???Does it work for heavy Elements?

Detailed Measurements ofComposition and Energy Spectra

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Secondary and Primary Cosmic Rays

“SECONDARY NUCLEI”, e.g. Li, Be, and B, are produced by spallation of “PRIMARY” parent nuclei in the ISM

~ 1 GeV/nucleon

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ISOTOPIC ABUNDANCES (measurements from ACE, 2001)

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RADIO-ACTIVE “CLOCK-NUCLEI”

measure containment life time of cosmic rays in Galaxy: about 15 M years at 1GeV/nucleon

[ 10 Be clock ]

measure time delay between nucleosynthesis

of primary nuclei and time of acceleration:

at least 105 years

[ 59Ni 59Co ]

SECONDARY/PRIMARY ABUNDANCE RATIOS vs. ENERGY (Data from ACE/CRIS and HEAO)

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DATA FROM SPACE: HEAO-3 (1990) AND CRN (1990)

DECREASE OF THE “L/M” ABUNDANCE RATIO:

Abundances of secondaryelements like Borondecrease with energyrelative to the abundancesof primary “parents” such ascarbon.[Juliusson, Meyer, Müller 1972]

The interstellar propagationpathlength Λ decreases withenergy (at least up to about100 GeV/n):

Λ E-0.6

(above 10 GeV/n)

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•Observer

GALAXY

Observed energy spectrum is power law E -2.7

Energy spectrum of particles injected by the source is different from observed spectrum:

With Λ E-0.6 hard source energy spectrum is required,in first order source power law E -2.1

* Source(whatever it is)

--- Λ(E) ---

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STOCHASTIC ACCELERATION IN STRONG SHOCKS IN SN REMNANTS:

Proposed 1977/78 by Axford et al; Bell; Blandford& Ostriker; and others.

Predicts hard source energy spectrum, about E-2

This is similar to what the measurements indicate!

This Process has now become the consensus model

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TeV GAMMA-RAY EMISSION FROM SHELL-TYPE SUPERNOVA REMNANTS (DATA: HESS 06)

RXJ1713.7-3946 Vela JuniorContours: ASCA 1-3 keV x-rays

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HISTORY

Baade and Zwicky 1934Supernova Explosions Fermi 1949

Distributed Acceleration But how does it work??? No magnetic Fields?

Efficiency???Does it work for heavy Elements?

Detailed Measurements ofComposition and Propagation

Bell and others, ~1980Stochastic Shock Acceleration in SNR

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SN shock acceleration is widely accepted, but questions remain:

The accelerator is expected to “run out of steam” at energies around Z x 1014 eV (Z = nuclear charge number).

New Measurements reaching this energy are needed!

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Galactic Cosmic Rays

ExtragalacticContributions?

THE EXPERIMENTAL CHALLENGE:

Below ~ 1010 eV/n :Solar modulation distortsenergies and spectra

Above the knee:No identification of individual particles

In between:Accurate measurements possible and necessary, but increasingly difficult at higher energies.

E-2.7

E -3.0

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DIRECT OBSERVATIONS :

Measured Quantities:

Charge Z (chemical identity) relatively easy

Mass M (isotopic species) extremely difficult

Energy E, or Lorentz factor E/mc2, or velocity v

Direction and trajectory through detector

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ENERGY MEASUREMENT

Detectors of 1 m2 area or more required Calorimeter no energy limit, but heavy

Magnet spectrometer up to 1000 GV; heavy and complex

Cherenkov counter gas counters up to several 100 GeV/amu

Relativistic rise of dE/dx in gases up to 1000 GeV/amu

Transition radiation det. up to 105 GeV/amu

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CRNP. Meyer,D. MüllerS. Swordy(1985)

“CRN”

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TRACER Detector System“Transition Radiation Array for Cosmic Energetic Radiation”

Scintillator 1Cherenkov 1

dE/dx Array

TRD4 Modules

Scintillator 2Cherenkov 2

2 m2 m

1.2 m

1600 proportional tubes, 2 cm dia, 200 cm long

Radiator

ICRC 2007 Merida

TRACER IS BIG: 5 m2 ster Currently the largest balloon-borne cosmic-ray detector

AND HEAVY: 5,000 lbs, 250 Watt, 1 Mbit/sec data

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Cherenkov

dE/dx

TRD

LORENTZ FACTOR γ

SIG

NA

L (

arb

. un

its)

ENERGY RESPONSE: Acrylic Cherenkov Counter (γ < 10) Specific Ionization in Gas (4 < γ < 1000) Transition Radiation Detector (γ > 400)

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ANTARCTICA 2003

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KIRUNA, SWEDEN 2006

TRACERTrajectory 2006

Complete circlearound pole not possible:Flight over Russianot permitted!

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CHARGE IDENTIFICATION

E

Z

Resolution (in charge units) O: 0.3 Fe: 0.5

Square Root ofScintillator and Cerenkov Signals

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• Radiators made from plastic fibers

• Previously used on CRN detector

• Calibrated at accelerators with

singly charged particles[L‘Heureux et al., 1990] Phys. Res. 295, 246, 1990]

TRD energy response measured in Xenon gas proportional counters

Lorentz factor γTRD response, arb. units

TR

D r

espo

nse,

arb

. un

its

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NEON NUCLEI 2003 Flight

SUB-RELATIVISTIC PARTICLES EXCLUDEDBY REQUIRINGCERENKOVIN SATURATION

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Previous Results from Space (HEAO-3 and CRN)

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Results from TRACER 2003

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BEST FIT POWER LAW INDEX FOR INDIVIDUAL ELEMENTS ABOVE 20 GeV/n

DATA FROM TRACER

E – 2.67

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PROPAGATION MODEL

AMBIENTCOSMIC RAYS,COMPONENT (i)

ESCAPE OR INTERACTION

COSMIC-RAYSOURCE

PRODUCTIONBY SPALLATION

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B/C RATIODECREASING WITH ENERGY

BUT LARGEUNCERTAINTIESBEYOND 100 GeV/n Λ(E) = b E-0.6 + Λ0

g/cm2

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PROPAGATION MODEL

Assume propagation pathlength Λe (E) = A E-0.6 + Λ 0

Fit data with three free parameters: power law at source, α residual pathlength, Λ0

abundance at source, qi

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Fitting results for the energy spectra of oxygen and iron

3σ

FIT FOR THE COMBINED SPECTRA OF ALL PRIMARY NUCLEI FROM O TO Fe

COMBINED FIT

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SUMMARY OF FITTING RESULTS

• All energy spectra can be simultaneously fit with a fairly soft source spectral index, α ≈ 2.35 to 2.45

• The residual pathlength Λ0 is not strongly constrained, possible values are Λ0 ≈ 0.1 to 0.5 g/cm2 , with larger values excluded by current L/M measurements.

• Relative source abundances of the elements are consistent with results from measurements at lower energies, and show similar correlations with FIP or volatility.

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Protons and Helium

Model Predictions for(α,Λ0) = (2.4, 0.3 g/cm2)

COSMIC-RAY ELECTRONS (conventional wisdom)

For high energies (>50 GeV),radiative energy losses dominant during propagation:

dE/dt = -k E2

Consequently, simple leaky box with power law source

energy spectrum E-α predicts observed spectrum E-(α+1).

q E-α = N(E)/TD + N(E)/TR

with TD diffusive life time

TR = 1/kE radiative life time

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Electrons (e++ e-): Differential Energy Spectrum, multiplied with E3

HEAT

Conventional Wisdom:Spectral shape corresponds to source spectrum E -2.3, like that of nuclei.

Data sets normalized at 10 GeV

Hal

oH

alo

Disk

λD(E)λD(E)

E increases

MORE REALISTIC DIFFUSION MODEL FOR ELECTRONS

Containment volume decreases with increasing energy E:

Diffusion coefficient D ~ E 0.6 Diffusion length λ(E) ~ E-0.4

Then observed energy spectrum N(E) ~ D-1/2 E-(α+0.5) = E-(α+0.8)

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Electrons (e++ e-): Differential Energy Spectrum, multiplied with E3

HEAT

Diffusion model: Spectral Shape corresponds to source ~ E -2.5

Data sets normalized at 10 GeV

Energy spectrum of Electrons ( e+ + e- )reported by ATIC (2008)

ELECTRON ENERGY SPECTRUM FROM FERMI/GLAST

ELECTRON ENERGY SPECTRUM FROM HESS (2009)

Interpretation of Electron Measurements

“Old Data”: Observed Spectral index 3.30+/- 0.06 source index α ≈ 2.50

ATIC Results: Spectral feature requires additional source contribution

Fermi Data: Observed spectral index 3.04 (+/- 0.05?) source index α ≈ 2.24 [note: Fermi authors propose α < 2.54]

Hess Results: Observed index 3.0 +/- 0.1 below 0.9 TeV source index α ≈ 2.2

CONCLUSIONS

• Energy spectra for the major primary cosmic-ray elements now are determined to energies above 1014 eV/particle. All energy spectra have similar shape.

• A simultaneous and self-consistent fit to all spectra is possible, but requires a fairly soft spectrum at the source, α ≈ 2.35 to 2.45.

• The energy dependence of the propagation pathlength is poorly constrained by the fit. More accurate measurements of the L/M ratio into the TeV/n region are needed and possible with current instrumentation.

• There are no striking discrepancies if the spectral fit for nuclei is applied to the current measurements of protons and electrons, but more work is required.