Lecture 14 Cosmic Rays

Transcription

1 Lecture 14 Cosmic Rays 1. Introduction and history 2. Locally observed properties 3. Interactions 4. Demodulation and ionization rate 5. Midplane interstellar pressure General Reference MS Longair, High Energy Astrophysics (especially Volume I Ch. 9)

2 1. Introduction DISCOVERY Part of the rise of modern physics: early radiation detectors (ionization chambers, electroscopes) showed a dark current in the absence of sources. Rutherford (1903): most comes from radioactivity Wulf (1910): dark current down by 2 at top of Eiffel Tower - could not be gamma rays Hess (1912): 5 km open-balloon flight showed an increase Hess & Kohlhörster (by 1914): balloon flights to 9 km... Studies of the variation with height, latitude & longitude confirmed the particle nature of cosmic rays (Millikan s name) originating above the Earth s atmosphere.

3 Role in Physics and Astrophysics Anderson discovered the positron in 1932 and shared the Nobel Prize with Hess in Until 1952, cosmic ray research was experimental particle physics. It led to many discoveries: muon, pion, and other particles. Even today the energy of the highest energy cosmic rays, > ev, is very much greater than available with accelerators. The extraterrestrial nature of cosmic rays might have early confronted astronomers, but this challenge could not be faced until mid 20 th century. Today cosmic rays are an important part of solar system and galactic astrophysics, including the ISM.

4 Remarkable Cosmic Ray Spectrum from ev Power law defined from ev steepens at the knee and recovers beyond the ankle. The puzzle of ultra highenergy CRs is that they can t be confined by the Galactic B-field, they can t be produced by SNe, and they can t come from very large distances because they interact with CMB photons. knee ankle GZK limit

5 2. Locally Observed Properties From observations above the atmosphere using balloons, rockets, satellites we know (Longair I Ch. 9) the following: 1. high degree of isotropy 2. power law spectra from ev (and higher)* 3. low-energy CRs excluded from solar system* 4. (mainly) solar abundances* 5. short lifetime in the Milky Way (20 Myr) c.f. detection of radioactive 10 Be (half-life 1.5 Myr) 6. significant pressure: ~10-12 dynes cm -2 *Illustrated below with figures

6 Cosmic Ray Spectra Simpson, ARNS Intensity vs. energy per nucleon from Mev/A. The units of intensity are: particles per (m 2 s MeV/nucleon). The proton slope is I p ( E) = ( E ) GeV 2.75 cm 2 s 1 sr 1 GeV 1

7 Cosmic Ray Abundances The excesses are largely due to spallation reactions of protons with abundant nuclei that can produce elements That ordinarily are produced in stars at low abundances.

8 The Electron Spectrum Although similar to protons, the electrons are even more reduced in intensity at low energies. The distribution is also affected by energy loss from synchrotron emission.

9 Most CRs have E ~ 1 GeV. They interact primarily with atomic electrons, exciting & ionizing atoms, as is well known from experiment (see figure) and from Bethe s theory. 3. Interactions Range-Energy Relation Curves vary as 5/3 power Energy per nucleon from MeV/A NB: The ordinate is the range multiplied by (Z 2 /A) for the projectile. The range of a 10-MeV proton is only ~ 1 gr cm -2

10 Nuclear and Magnetic Interactions At GeV energies, the nuclear cross section is ~ 10 mb, equivalent to ~ 100 gr cm -2, thus nuclear are less important than electronic interactions. But scattering from bent or kinked magnetic field lines can be more important. It arises from the inability of a charged particle to continue spiraling around a magnetic field when the fields vary rapidly in space. Another important CR interaction is with Alfven waves. These magnetic processes are especially important where the interstellar turbulence is MHD in nature. The affect the transport of CRs through the galaxy.

11 4. Demodulation With I(E) decreasing rapidly with E, it is important to understand the low-energy behavior. The figure shows spectra at three levels of solar activity, indicating that the Sun itself changes the CR intensity. Correcting the observed spectra for solar system effects, or demodulation, is required to deduce the CR intensity in the local ISM.

12 Effects of the Earth s Magnetic Field The effects of the Earth s magnetic field have been studied extensively for more than 100 years and are reasonably well understood. Satellite observations of CRs extend beyond this region.

13 Interfaces in the Heliosphere Bow shock c.f. interstellar wind Terminal shock c.f. solar wind Suess, Rev Geophys Deducing the demodulated (or true) CR intensity makes use of satellite observations as a function of heliocentric distance and transport theory. The satellite observations are now approaching the crucial terminal shock region.

14 Satellite Locations vs. Time

15 Cosmic Ray Ionization Rate Given a demodulated CR intensity, the ionization rate can be calculated by integration with the ionization cross section. H, He, and H 2 are the most important targets. A basic fact is that 37 ev is needed to make an ion pair, so a 2 MeV proton makes about 54,000 ions. The CR ionization rate per proton is ζ = f f ζ f CR sec 5 3 heavy f sec heavy Integrating down to 2 MeV, Spitzer & Tomasko (ApJ ) found: ζ p s 1 and CR p With a 1975 demodulation, the rate increases to ς CR = 5x10-17 s -1 ζ s 1

16 Cosmic Ray Ionization Rate Update Webber (ApJ ) used satellite observations out to 42 AU (c. 1987). Repeating the demodulation calculations, the ionization rate down to 10 MeV/nucleon is ζ CR = (3 4) s 1 Going down to 2 MeV would increase this by ~ 50%.

17 5. Midplane Pressure References: McKee in Evolution of the ISM (ASP 1990), p. 3 Boulares & Cox, ApJ (BC) It is generally assumed that, despite itremendous dynamic activity, the Milky Way is in hydrostatic equilibrium, its stability guaranteed by a large midplane pressure (from several components) and by a large halo. For simplicity, we assume that the ISM is vertically stratified and satisfies the usual equation dp dz = ρ( z) g( z) where g(z) is the gravitational acceleration, mainly due to stars.

18 The solution Hydrostatic Equilibrium (cont d) z p( z) = dz' ρ( z') g( z') requires a knowledge of both ρ(z) and g(z). We know that the former is very uncertain, but the latter is also uncertain by 25-35%. McKee and BC make estimates of the contributions to ρ(z) of the observed phases to estimate p(0). Then they see whether the result is consistent with what we know about the pressure contributions: p = ( p + p + p th + pturb + pram) For example, the main gas contributions included by BC are given on the next page. mag CR

19 Simplified Boulare & Cox Model Phase Density (cm -3 ) Scale height (pc) H CNM WNM WIM The total half-column of gas is about 5x10 20 cm -2, which corresponds to a (half) surface density of 6 M Sun pc -2, cool and the warm contributing roughly equal amounts.

20 Boulares-Cox Midplane Pressure Estimate 12-3 p(0) (3.9 ± 0.6) 10 erg cm If this came from purely gas kinetic (thermal) pressure, the nt product would be nt = p(0) / k B 28,000 cm The thermal pressure determined from optical & UV absorption lines is ~ 3,000 cm -3 K, ~ 10 times too small. From this and the precious lecture, we get p mag p CR 10 which together account for half of the total. The other 40% might be turbulent pressure, a more dynamic ram pressure, or larger magnetic and/or CR pressure. 12 erg 3 cm K -3