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The Theory/Observation connection lecture 2 perturbations

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Configuration space descriptionConfiguration space descriptionConfiguration space descriptionConfiguration space description

Lecture outline

- Describing perturbations
- correlation function
- power spectrum
- Perturbations from Inflation
- The evolution on perturbations before matter dominated epoch
- Effects from matter-radiation equality
- Effects from baryonic material
- evolution during matter domination
- linear growth
- linear vs non-linear structure (introduction)

Perturbation statistics: correlation function

overdensity

field

definition of

correlation function

from statistical homogeneity

from statistical isotropy

can estimate correlation function using galaxy (DD) and random (RR) pair counts at separations ~r

Perturbation statistics: power spectrum

definition of

power spectrum

power spectrum is the Fourier analogue of the correlation function

sometimes written in

dimensionless form

The importance of 2-pt statistics

Because the central limit theorem implies that a density distribution is asymptotically Gaussian in the limit where the density results from the average of many independent processes; and a Gaussian is completely characterised by its mean (overdensity=0) and variance (given by either the correlation function or the power spectrum)

Correlation function vs Power Spectrum

The power spectrum and correlation function contain the same information; accurate measurement of each give the same constraints on cosmological models.

Both power spectrum and correlation function can be measured relatively easily (and with amazing complexity)

The power spectrum has the advantage that different modes are uncorrelated (as a consequence of statistical homogeneity).

Models tend to focus on the power spectrum, so it is common for observations to do the same ...

Phases of (linear) perturbation evolution

Inflation

Matter/Dark energy

domination

Transfer function

Non-linear

(next lecture)

linear

Why is there structure?

Inflation (a period of rapid growth of the early

Universe driven by a scalar field) was postulated to solve some serious problems with “standard” cosmology:

- why do causally disconnected regions appear to have the same properties? – they were connected in the past (discussed in last lecture)
- why is the energy density of the Universe close to critical density? – driven there by inflation
- what are the seeds of present-day structure? -Quantum fluctuations in the matter density are increased to significant levels

Driving inflation

Every elementary particle (e.g. electron, neutrino, quarks, photons) is associated with a field. Simplest fields are scalar (e.g. Higgs field), and a similar field (x,t) could drive inflation.

Energy-momentum tensor for (x,t)

For homogeneous (part of) (x,t)

Acceleration requires field with

Driving inflation

Acceleration requires field with more potential energy than kinetic energy

Energy density is constant, so Einstein’s equation gives that the evolution of a is

Gives exponential growth

Problems: only way to escape growth is quantum tunneling, but has been shown not to work.

Instead, think of particle slowly rolling down potential. Close to, but not perfectly stationary.

Inflation: perturbations

Perturbations in the FRW metric

Scalar potentials usually give rise to

so can interchange these

Spatial distribution of fluctuations can be written as a function of the power spectrum

During exponential growth, there is no preferred scale, so (Gaussian) quantum fluctuations give rise to fluctuations in the metric with

with n=1

Inflation: perturbations

Poisson equation translates between perturbations in the gravitational potential and the overdensity

In Fourier space

So matter power spectrum has the form

Because quantum fluctuations are Gaussian distributed, so are resulting matter fluctuations, which form a Gaussian Random Field

Phases of perturbation evolution

Inflation

Matter/Dark energy

domination

Transfer function

linear

Non-linear

Jeans length

- After inflation, the evolution of density fluctuations depends on the scale and composition of the matter (CDM, baryons, neutrinos, etc.)
- An important scale is the Jeans Length which is the scale of fluctuation where pressure support equals gravitational collapse,

where csis the sound speed of the matter, andis the density of matter.

“F=ma” for perturbation growth

depends on Jeans scale

Transfer function evolution

in radiation dominated Universe, pressure support means that small perturbations cannot collapse (large Jeans scale). Jeans scale changes with time, leading to smooth turn-over of matter power spectrum. Cut-off dependent on matter density times the Hubble parameterWmh.

large scale

perturbation

gravitational potential k3/2F

small scale

scale factor of Universe a

The power spectrum turn-over

Can give a measurement of the matter density from galaxy surveys

It is hard to disentangle this shape change from galaxy bias and non-linear effects

varying the matter density

times the Hubble constant

The effect of neutrinos

The existence of massive neutrinos can also introduce a suppression of T(k) on small scales relative to their Jeans length. Degenerate with the suppression caused by radiation epoch. Position depend on neutrino-mass equality scale.

kP(k)

Configuration space description

Wm=0.3,Wv=0.7, h=0.7,Wb/Wm=0.16

position-space description: Bashinsky & Bertschinger

astro-ph/0012153 & astro-ph/02022153 plots by Dan Eisenstein

Configuration space description

Wm=0.3,Wv=0.7, h=0.7,Wb/Wm=0.16

position-space description: Bashinsky & Bertschinger

astro-ph/0012153 & astro-ph/02022153 plots by Dan Eisenstein

Configuration space description

Wm=0.3,Wv=0.7, h=0.7,Wb/Wm=0.16

position-space description: Bashinsky & Bertschinger

astro-ph/0012153 & astro-ph/02022153 plots by Dan Eisenstein

Wm=0.3,Wv=0.7, h=0.7,Wb/Wm=0.16

position-space description: Bashinsky & Bertschinger

astro-ph/0012153 & astro-ph/02022153 plots by Dan Eisenstein

Wm=0.3,Wv=0.7, h=0.7,Wb/Wm=0.16

position-space description: Bashinsky & Bertschinger

astro-ph/0012153 & astro-ph/02022153 plots by Dan Eisenstein

Wm=0.3,Wv=0.7, h=0.7,Wb/Wm=0.16

position-space description: Bashinsky & Bertschinger

astro-ph/0012153 & astro-ph/02022153 plots by Dan Eisenstein

Wm=0.3,Wv=0.7, h=0.7,Wb/Wm=0.16

position-space description: Bashinsky & Bertschinger

astro-ph/0012153 & astro-ph/02022153 plots by Dan Eisenstein

Baryon Oscillations in the power spectrum

“Wavelength” of baryonic acoustic oscillations is determined by the comoving sound horizon at recombination

At early times can ignore dark energy, so comoving sound horizon is given by

varying the

baryon fraction

Sound speed cs

Gives the comoving sound horizon ~110h-1Mpc, and BAO wavelength 0.06hMpc-1

The matter power spectrum

Overall shape of matter power spectrum is given by

Current best estimate of the galaxy power spectrum from SDSS: no sign of turn-over yet

Phases of perturbation evolution

Inflation

Matter/Dark energy

domination

Transfer function

linear

Non-linear

Spherical perturbation leading to linear growth

cosmology equation

homogeneous dark energy means that this term depends on scale factor of background

“perfectly” clustering dark energy – replace a with ap

Consider homogeneous spherical perturbation – evolution is “same” as “mini-universe”

Overdense perturbation

Radius ap

Background

Radius a

Spherical perturbation leading to linear growth

cosmology equation

definition of d

to first order in perturbation radius

(linear approximation)

gives

can also be derived using the Jeans equation

only has this form if the dark energy does not cluster – derivation of equation relies on cancellation in dark energy terms in perturbation and background

Linear growth: EdS model

For flat matter dominated model, this has solution

Remember that the gravitational potential and the overdensity are related by Poisson’s equation

Then the potential is constant: there is a delicate balance between structure growth and expansion

Not true if dark energy or neutrinos

Linear growth: general models

For general models, denote linear growth parameter (solution to this differential equation)

For lambda models, can use the approximation of Carroll, Press & Turner (1992)

For general dark energy models, need to solve the differential equation numerically

Linear vs Non-linear behaviour

z=0

non-linear

evolution

z=1

z=2

z=3

linear

growth

z=4

z=0

z=5

large scale power

is lost as fluctuations

move to smaller scales

z=1

z=2

z=3

z=4

z=5

P(k) calculated from Smith et al. 2003, MNRAS, 341,1311 fitting formulae

Further reading

- Dodelson, “Modern Cosmology”, Academic Press
- Peacock, “Cosmological Physics”, Cambridge University Press
- Liddle & Lyth, “Cosmological Inflation and Large-Scale Structure”, Cambridge University Press
- Eisenstein et al. 2006, astro-ph/0604361 (configuration space description of perturbation evolution)
- Percival 2005, astro-ph/0508156 (linear growth in general dark energy models)

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