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High Energy Astrophysics > 10: COSMIC RAY ELECTRONS AND PROTONS > Flashcards

10: COSMIC RAY ELECTRONS AND PROTONS Flashcards

1
Q

Why are direct observations of CR electrons only possible at >10 GeV? What is observed?

A

At lower energies, they lose flux to the magnetosphere and solar wind B-field, so can’t be measured reliable.

Solar modulation is observed - fewer CRs detected when the sun is more active.

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2
Q

Between ~10 GeV and 10^3 GeV, direct measurements give what and what is p equal to?

A

Power-law spectrum. P is equal to 3.3.

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3
Q

Why do we see synchrotron radiation from approaching CR electrons?

A

Because of the Galactic B field.

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4
Q

What can we use the slope alpha of a synchrotron spectrum for?

A

To infer the slope p of the energy spectrum of emitting electrons.

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5
Q

Measurements of the radio emission from the Galactic disk are possible between roughly 10 MHz and 2 GHz. Why doesn’t it work at higher/lower frequencies?

A

The signal is too faint at higher frequencies and at lower frequencies the radiation is lost to bremsstrahlung absorption.

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6
Q

What range do synchrotron observations show electrons in?

A

The range 0.5 - 8 GeV. They have p ~ 1.8 - 2.6.

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7
Q

Below 0.1 GeV, why do we estimate the CR spectral slope from the Galactic gamma ray spectrum?

A

Electron bremsstrahlung on the ISM dominates the gamma-ray emission.

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8
Q

Why measurement is produced by the gamma-ray spectrum?

A

p ~ 1.6

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9
Q

What is the slope and energy range (GeV) for:
1. gamma rays
2. radio synchrotron
3. direct measurements

A
  1. E^-1.6, <0.1 GeV
  2. E^-1.8 to E^-2.6, 0.5 - 8 GeV
  3. E^-3.3, > 10
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10
Q

What are the four energy-loss processes and their energy dependence?

A
  1. Ionisation - Energy-loss rate is constant with energy
  2. Bremsstrahlung - Energy is proportional to loss rate.
  3. Synchrotron - Loss rate proportional to E^2
  4. Inverse compton - Loss rate proportional to E^2
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11
Q

Define ionisation.

A

Electrons kick electrons out of atoms by electrostatic repulsion. Protons have a weak dependence on energy.

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12
Q

Define Bremsstrahlung.

A

Coloumb-force acceleration causing EM radiation. Stronger if atoms ionised. Considering energy loss rate from a single electron rather than a population.

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13
Q

Define synchrotron.

A

B-field-force acceleration causing EM radiation.

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14
Q

Define Inverse Compton.

A

Photons in the radiation field of energy density u_rad upscattered at the expense of electron energy.

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15
Q

Where do ionisation, synchrotron, and IC losses dominate?

A

Ionisation - low energies
Sync/IC - high energies

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16
Q

Why can’t high-energy electrons have cosmological origin?

A

Because of energy loss to iC on CMB. They would lose all their energy to IC emission before travelling any great distance.

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17
Q

What is the value of gamma, the energy-loss lifetime and the max distance of high energy CR electrons?

A

gamma ~ 2 x 10^5
energy-loss lifetime on CMB ~ 10^7 years
max distance at the speed of light ~ 3Mpc

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18
Q

What is the peak energy of CMB photon?

A

7 x 10^-10 MeV

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19
Q

What does n(E) represent?

A

number of electrons per unit energy per unit volume

20
Q

What does Q(E) represent?

A

number of new particles created per unit energy per unit time (the source term)

21
Q

What happens in steady state dn/dt=0?

A

At all energies, energy losses balanced by new particles.

22
Q

What does n(E)(−dE/dt) represent?

A

Number of e− losing energy dE in time dt per unit volume

23
Q

What does the ∂/∂E term describe?

A

The difference between the number of particles moving into and out of
energy range E → E + dE due to energy loss.

24
Q

For the residence time in the Galaxy equation, what are the integral limits and what do the left and right hand side of the equation represent?

A

n —> 0 as E —> infinity
LHS is number e− cooling to < E per unit time & volume
RHS is the number new e− > E produced by a source per unit time & volume

25
Q

What are the important energy loss terms?

A

A_1 - ionisation
A_2E - Bremss
A_3E^2 - syn, IC

26
Q

If the energy losses are dominated by ionisation, Bremsstrahlung or Synchrotron/IC, how does the proportionality and gradient relationship with respect to the source change?

A

Ionisation: n ∝ E^(−(p−1)) - Flatter than source by 1
Brem: n ∝ E^(−p) - Unchanged
Sync/IC: n ∝ E^(-(p+1)) - Steeper than source by 1

27
Q

What dominates at the measurement method of low-energy gamma rays, what is the energy to the power of, what is the value of E roughly, and what is the loss time?

A

Ionisation losses.
E is to the power of -1.6
E < or = to 0.1 GeV
10^7 years

28
Q

What dominates at the measurement method of direct measurement, what is the energy to the power of, what is the value of E roughly, and what is the loss time?

A

Syn, iC losses.
E is to the power of -3.3
E greater or = to 10 GeV
10^7 years

29
Q

What dominates at the radio synchrotron observations, what is the energy to the power of, what is the value of E roughly, and what is the loss time?

A

Escape of the electrons before energy loss. An unmodified source spectrum is seen.
E to the power of -1.8 to -2.6
E is roughly 1 - 10 GeV
Electrons escape before loss

30
Q

What happens when p = 2.4?

A

Breaks in the observed CR electron spectrum is observed due to electrons escaping the Galaxy in 10^7 years, but they first suffer ionisation losses at low energies and synch/iC at high energy.

31
Q

What would happen if the residence time was not 10^7 years?

A

The spectral breaks wouldn’t be where observed.

32
Q

What is the magnetic field in the SNR?

A

2.7 x 10exp(-8) T

33
Q

What percentage of electrons and positrons are primary e-?

A

80%

34
Q

What percentage of electrons and positrons are secondary e-?

A

10%

35
Q

Why is the positron spectrum steeper than the electron spectrum and why?

A

Because the positrons are secondary particles. The number of high-energy positrons drops off more quickly with increasing energy than the number of electrons.

If positrons are secondary, since secondary particles must be produced by a higher energy primary, it’s harder to produce very high energy secondaries.

36
Q

Why can’t cosmic ray proton spectra be observed below E < 1 Gev?

A

They suffer solar modulation.

37
Q

How do CR protons lose energy?

A

Through strong interactions - inelastic collisions with protons in the ISM produce mainly pions.

38
Q

What do charged pions decay to?

A

Secondary CR electrons and positrons.

39
Q

What is the neutral pion decay time?

A

0.84 x 10^-16 s

40
Q

What is the 99% of neutral pion decays?

A

π^0 → γ + γ

41
Q

What is the rest mass energy of the pion?

A

135 MeV

42
Q

How much does each gamma get from the pion interaction in the rest frame?

A

67.5 MeV

43
Q

In lab frame we see a spread of gamma-ray energies. How is this depicted in the local gamma-ray spectrum? What else is seen?

A

As a hump.
Soft gamma-ray bremsstrahlung emission that is used to infer the CR electron spectrum at low energies.

44
Q

What peaks towards the Galactic centre?

A

The gamma-ray emission > 100 MeV along the Galactic plane.

45
Q

Prediction with CR density, nCR = constant, using ISM measurements (HI, CO as tracer of H2) gave
insufficient contrast between the Galactic centre and anti-centre directions in the early SAS II data. What was a better fit? What did this suggest?

A

Much better fit if nCR ∝ nISM; i.e. emissivity
J(γ) ∝ nCRnISM ∝ n_ISM^2
Galactic origin for the CR protons

46
Q

Why do observations not suggest CR protons are extragalactic?

A

We would find them to be uniformly distributed through the Galaxy, but this assumption gives a prediction for the distribution of γ-rays produced by those CR protons that is too smooth compared with observations. Instead, if the CR protons are produced in the Galaxy at locations that trace the ISM, then the predicted γ-ray distribution is a much better match to observations.

High Energy Astrophysics (11 decks)

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