![]() | Annu. Rev. Astron. Astrophys. 1994. 32:
319-70 Copyright © 1994 by Annual Reviews. All rights reserved |
3.2. Power Spectra on Smaller Scales
Although on large angular scales the transfer function is
T(k) 1,
there is significant structure on small to intermediate scales.
Generally the Sachs-Wolfe effect dominates the power spectrum on
scales larger than the horizon size at last scattering and the power
spectrum can be taken to be a power law. On smaller scales, however,
causal interactions become important and the spectrum is modified.
For a given cosmological model, the shape of the fluctuation
spectrum is fixed, and depends on the primordial spectrum
(e.g. kn)
and its evolution as the waves enter the horizon. This makes the
spectrum at smaller scales dependent on
0,
B,
H0, and the dark
matter. Given a cosmological model, however, both the radiation and
matter power spectra are well defined. For a fixed
0, the
radiation power spectrum depends only on the type of dark matter at the
high k end.
The calculation proceeds by considering perturbations in the photon distribution function, f (x, q, t), which can be written in terms of
![]() | (22) |
where
is the
Planck function and q is the comoving photon
momentum.
is the
total energy, or brightness, perturbation, which is also
4
T / T
for a uniform shift in temperature. Liouville's theorem
tells us that the total phase space density is conserved for
collisionless particles, i.e. the distribution function is constant
along particle paths: Df /Dt = 0. Source terms
("collisions") add on
the right hand side; the important sources are baryon velocities and
photon density perturbations coupled through Thomson scattering. The
Boltzmann equation (or equation of radiative transfer, written here in
the synchronous gauge) for
is
![]() | (23) |
(see Peebles & Yu
1970,
Hu et al 1994,
Dodelson & Jubas
1994),
where the dot denotes differentiation with respect to conformal time
(=
dt / a),
i
are the direction cosines defined by
,
ne is the number
density of free electrons, a(t) is the cosmological
scale-factor, vB is the baryon peculiar velocity,
0 is the
isotropic part of
,
hij is
the metric perturbation, and isotropic scattering has been
assumed. The Boltzmann equation is generally Fourier transformed,
since in the linear approximation the different k-modes evolve
independently, which makes the calculation tractable. It is also
assumed that all fluctuations are still in the linear regime, which is
a reasonable approximation for the relevant scales at the scattering
epoch. Since the effects of the radiation on the matter cannot be
ignored (except for late reionization), the equation of motion for the
baryons needs to be solved simultaneously; this is the continuity
equation for matter evolving freely, but also takes into account the
Compton drag at early times.
The radiation and baryons evolve as a coupled fluid at early times,
but need to be followed more accurately as the Universe recombines,
and eventually the baryons decouple from the photons entirely. After
this point, the photons can be assumed to free-stream to the observer,
and the subsequent behavior of the anisotropies is often treated
analytically. Dark matter evolves collisionlessly throughout. although
its gravity feeds back into the baryon and photon evolution. Detailed
calculations along these lines have been carried out for many
different cosmological models, e.g. the work of
Peebles & Yu (1970),
Wilson & Silk (1981),
Bond & Efstathiou
(1984,
1987),
Vittorio & Silk (1984,
1992),
Holtzman (1989),
Sugiyama & Gouda
(1992),
Dodelson & Jubas
(1993a),
Stompor (1993), and
Crittenden et al (1993a,
b)
among others. The Boltzmann equation is usually solved by expanding the
s
in Legendre polynomials up to some high enough
and numerically
integrating the coupled equations.
We show in Figure 2 the power spectrum for CDM
models as a function of wavenumber k, for a range of
B
consistent with Big Bang Nucleosynthesis (BBN)
(Krauss & Romanelli
1990,
Walker et al 1991,
Smith et al 1993).
To calculate the
C
from
these power spectra
approximately, one integrates the power spectrum with measure
j
2(k
0)
dk / k. Each
C
is thus
in effect an average of the power spectrum around
k
H0/2. We
have also included an xe = 1 reionized
BDM power spectrum (arbitrarily normalized) on the figure for comparison.
![]() |
Figure 2. Power spectrum for "standard" CDM
models (h = 1/2,
|
The plateau in Figure 2 at low k is the
contribution from the
Sacks-Wolfe effect (or gravitational redshift), which damps on the
scale of the scattering surface thickness. The bumps and wiggles
reflect the phase of the oscillating baryons and photons when
recombination occurs, i.e. the number of oscillations a given mode has
undergone in the time between entering the horizon and the switch-off
of radiation pressure when the matter becomes neutral. These so-called
"Doppler" peaks (physically they come from a combination of v and
sources which are difficult to separate) rise at
k ~ 0.01 Mpc, the
size of the horizon at the time of last scattering. In
-space the
rise to the first Doppler peak is quite gradual, which can lead an
effective n > 1 even on relatively large scales, e.g.
neff
1.15 on COBE scales
(Bond 1994).
The height of the Doppler peaks is dependent
on the number of scatterers or
B
h2 (scaling roughly as
B1/3, which
would be the dependence for an experiment like MAX; see also
Fukugita et al 1990).
The second and third peaks are harmonics of the first
which is the fluctuation that underwent half an oscillation since it
entered the horizon. The amplitude of these oscillations reflects the
amount of growth before the perturbation enters the photon-baryon
Jeans scale, which depends on
B
h2. There is a further dependence on h through
zeq. There is also an effect due to the baryons falling
into the dark matter potential wells after recombination, although this
happens largely after the photons have been scattered. CDM
fluctuations that first entered the horizon during radiation
domination suffer growth suppression, which partly accounts for the
differing heights of the peaks. Note, also, the exponential cutoff in
the power spectrum at large k, due to the effects of the thickness of
the last scattering surface, as well as the familiar damping
(Silk 1968)
of baryon and photon fluctuations prior to decoupling. The
damping scale is set by the thickness
z (see
Appendix A).
For a reionized model (see Section 5.1)
the visibility function is
centered around some z* and has width
z ~
z*, so that the two
important scales above are almost the same. The effect is that there
is only one Doppler peak, which is at the scale of the horizon at
z*
(i.e. smaller k or
),
with damping for higher wavenumbers. There are
no extra bumps and wiggles, since at the time of last scattering the
photons and baryons were no longer oscillating.