The Diffuse Interstellar Medium ISM

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• The Diffuse Interstellar Medium (ISM)
• Lecture Topics
• the 21-cm line in emission and absorption
• heating-cooling balance and the phases of the
  ISM
• turbulence and structure in the ISM
• magnetic fields, polarization and Faraday
  tomography

John M. Dickey University of Tasmania February
2008
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Other tracers of the diffuse neutral medium
(other than the l21-cm line of H)

• Optical absorption lines (Na, Ca, Ti, )
• Mid-infrared and far-infrared fine structure
  lines
• optical obscuration (but dominated by denser
  clouds)
• molecular line emission (but few molecular
  species
• trace low densities well)
• carbon recombination lines at low radio
  frequencies
• diffuse Ha emission (warm ionized medium)
• pulsar dispersion and scattering (WIM HII
  regions)

Tracers of the diffuse ionized medium
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luminosity of the Milky Way disk
COBE results taken from Bennett et al. 1994, Ap.
J. 434, 587.
Fine structure lines are transitions between very
closely spaced electronic energy levels in atoms.
Typically these will be ionized away in HII
regions, so these lines select the diffuse,
neutral medium (dominated by PDRs).
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Note that both CI and CII are very powerful
lines. These are dominant coolants of the
interstellar medium. The molecular gas cools
primarily by molecular emission. Each of the
main CO lines (115, 230, 346, 461, 576, 691,
GHz) carries about 105 Lsun of energy out of the
Galaxy.
The contrast between CI and CII shows where the
radiation field is strong in the uv between 11.26
eV and 13.6 eV. Note that He, O, N, and many
other common atoms in the interstellar medium
have ionization potentials greater than 13.6 eV
(hydrogen). A large fraction of the interstellar
volume may be filled with CII, and CI may be
present only in higher density regions
(clouds). The situation for Na and Ca is
similar to C (a mixture of I and II).
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Thermal equilibrium (if it exists) requires
balance between the rates of heating and cooling,
indicated by G and L, respectively. These each
have units of watts m-3 or ergs cm-3 sec-1. They
generally have many different processes
contributing, and they typically have different
functional dependence on density and temperature.
For example, for the 158 mm CII line, the
cooling rate is
step function in T
where nH is the density in cm-3 and dC is the
carbon depletion (out of the gas phase onto
grains). Cooling by excitation of hydrogen
electronic transitions only becomes significant
when T gt 104 K. The cooling rate in this case
(for collisions with electrons) is
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Generally cooling is a collisional process, so it
goes as density squared
where f is some function of T including steps for
each fine structure line, and other cooling
processes. Heating processes generally involve
interaction of matter (atoms, molecules or grains
of dust) with the radiation field (or possibly
the magnetic field). These may have weak
dependence on temperature, but roughly they go
only as density to the first power
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A classical figure of the cooling function, f(T)
from Dalgarno and McCray, 1972, ARAA 10, 375.
log f(T)
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The heating rate due to photo-electron ejection
from grains is significant in many interstellar
environments. Here is an early analysis from
Draine 1978. The discovery of very small grains
(PAHs) increased the estimates of the cooling
rate by this process greatly.
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In the simple case that L a n2 f (T) and G a n,
then equilibrium requires something like
G L if the density mirrors the cooling
curve i.e. log n constant - log f
L gt G
log n
G gt L
log n0
log f(T)
2
3
4
5
log T
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L gt G
Swap the x and y axes.
log n
log n0
G gt L
2
3
4
5
log T
5
L gt G
4
G gt L
3
log T
2
log n0
log n
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Change variables on the y axis, from Temperature
to pressure, P log P log nT log n log
T This is like adding a rising line with slope 1
to the equilibrium curve.
L gt G
G gt L
L gt G
log n
log T log n log nT log P
G gt L
log n0
log n
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lines of constant T
T 6000 K warm phase
L gt G
log P
T 60 K cool phase
G gt L
negative slope makes this segment unstable
G L
log n
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from Wolfire et al. 1995 Ap. J. 443, 152
log P
log n
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The equilibrium curve changes with the dust to
gas ratio and with the heavy element
abundance. We would expect variation in fc, the
cool phase fraction of the total gas density, in
different galaxies and in the Milky Way outer
disk.
Wolfire et al. 1995
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Remember this table from lecture 1
It seems that lower heavy element abundance
causes lower cool phase fraction in galaxies
overall. The same thing appears to be true in the
outer Milky Way, far beyond the solar circle
(Strasser 2006, Ph.D. thesis).
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Going up out of the disk into the halo of the
Milky Way, what happens?
z
Fz
Fz
force toward the plane
slope
z
300 pc
Oort 1932, BAIN 6, 249.
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In an isothermal atmosphere the pressure drops
with altitude to balance the overburden, i.e. the
gravitational force of the mass of the gas above.
This gives
The second equality is true only within about 300
pc of mid-plane, where the gravitational
potential is parabolic with
At higher z the density distribution becomes an
exponential.
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Note that each phase has its own rms velocity
distribution, for the cool phase it is the
cloud-cloud motions that set the scale height,
not the microscopic thermal motions.
The pressure must decrease with z, so that at
some height above (or below) the plane the cool
phase can no longer exist in equilibrium.
cool phase
high pressure
log P
medium pressure
low pressure (z gt 500 pc ?)
warm phase
log n
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So we expect the mixture of phases to reflect
the rms pressure. In a dynamic ISM (with
supernova shocks, superbubbles, etc.) the
pressure fluctuates with time at every point. So
thermal equilibrium analysis will not explain
all features of the gas distribution. But after
a cooling time, the gas will try to return to
heating-cooling equilibrium.
is the cooling time.
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Dickey and Lockman, 1990, ARAA 28, 215.
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