2. The Physics of the cosmic microwave background Bonn, August 31, 2005 Ruth Durrer Départment de physique théorique, Université de Genève

3. Contents Introduction Linear perturbation theory - perturbation varibles, gauge invariance - Einstein’s equations - conservation & matter equations - simple models, adiabatic perturbations - lightlike geodesics - polarisation Power spectrum Observations Parameter estimation - parameter dependence of CMB anisotropies and LSS - reionisation - degeneracies Conlusions

4. The cosmic micro wave background, CMB After recombination (T ~ 3000K, t~3x10 5 years) the photons propagate freely, simply redshifted due to the expansion of the universe The spectrum of the CMB is a ‘perfect’ Planck spectrum: -4  < 10 -4 y < 10 -5

5. CMB anisotropies WMAP (2003) COBE (1992)

6. The CMB has small fluctuations,  T/T » a few £ 10 -5. As we shall see they reflect roughly the amplitude of the gravitational potential. => CMB anisotropies can be treated with linear perturbation theory. The basic idea is, that structure grew out of small initial fluctuations by gravitational instability. => At least the beginning of their evolution can be treated with linear perturbation theory. As we shall see, the gravitational potential does not grow within linear perturbation theory. Hence initial fluctuations with an amplitude of » a few £ 10 -5 are needed. During a phase of inflationary expansion of the universe such fluctuations emerge out of the quantum fluctuations of the inflation and the gravitational field.

7. metric perturbations Decomposition into scalar, vector and tensor components Linear cosmological perturbation theory

8. Perturbations of the energy momentum tensor Density and velocity stress tensor

9. Gauge invariance Linear perturbations change under linearized coordinate transformations, but physical effects are independent of them. It is thus useful to express the equations in terms of gauge-invariant combinations. These usually also have a simple physical meaning.  is the analog of the Newtonian potential. In simple cases  = . In longitudinal gauge, the metric perturbations are given by h  (long) = -2  d  2 -2  ij dx i dx j Gauge invariant metric fluctuations (the Bardeen potentials) _

10. The Weyl tensor The Weyl tensor of a Friedman universe vanishes. Its perturbation it therefore a gauge invariant quantity. For scalar perturbations, its ‘magnetic part’ vanishes and the electric part is given by E ij = C  ij u  u = ½[  i  j (  +  ) -1/3  (  +  )]

11. Velocity and density perturbations Gauge invariant variables for perturbations of the energy momentum tensor The anisotropic stress potential  The entropy perturbation w=p/  c 2 s =p’/  ’

12. Einstein equations constraints + dynamical _ Conservation equations +

13. matter The D 1 -mode is singular, the D 2 -mode is the adiabatic mode In a mixed matter/radiation model there is a second regular mode, the isocurvature mode On super horizon scales, x<1,  is constant On sub horizon scales, D g and V oscillate while  oscillates and decays like 1/x 2 in a radiation universe. Simple solutions and consequences radiation x=c s k 

14. radiation in a matter dominated background with Purely adiabatic fluctuations, D gr = 4/3 D m Simple solutions and consequences (cont.)

15. ++ RD ‘90 lightlike geodesics From the surface of last scattering into our antennas the CMB photons travel along geodesics. By integrating the geodesic equation, we obtain the change of energy in a given direction n: E f /E i = (n. u) f /(n. u) i = [T f /T i ](1+  T f /T f -  T i /T i ) This corresponds to a temperature variation. In first order perturbation theory one finds for scalar perturbations acoustic oscillations Doppler term gravitat. potentiel (Sachs Wolfe) integrated Sachs Wolfe ISW

16. Polarisation Thomson scattering depends on polarisation: a quadrupole anisotropy of the incoming wave generates linear polarisation of the outgoing wave.

17. Polarisation can be described by the Stokes parameters, but they depend on the choice of the coordinate system. The (complex) amplitude  i e i of the 2-component electric field defines the spin 2 intensity A ij =  i *  j which can be written in terms of Pauli matrices as Q § iU are the m = § 2 spin eigenstates, which are expanded in spin 2 spherical harmonics. Their real and imaginary parts are called the ‘electric’ and ‘magnetic’ polarisations (Seljak & Zaldarriaga, 97, Kamionkowski et al. ’97, Hu & White ’97)

18. E is parity even while B is odd. E describes gradient fields on the sphere (generated by scalar as well as tensor modes), while B describes the rotational component of the polarisation field (generated only by tensor or vector modes). E-polarisation (generated by scalar and tensor modes) B-polarisation (generated only by the tensor mode) Due to their parity, T and B are not correlated while T and E are

19. An additional effect on CMB fluctuations is Silk damping: on small scales, of the order of the size of the mean free path of CMB photons, fluctuations are damped due to free streaming: photons stream out of over-densities into under-densities. To compute the effects of Silk damping and polarisation we have to solve the Boltzmann equation for the Stokes parameters of the CMB radiation. This is usually done with a standard, publicly available code like CMBfast (Seljak & Zaldarriaga), CAMBcode (Bridle & Lewis) or CMBeasy (Doran).

20. Reionization The absence of the so called Gunn-Peterson trough in quasar spectra tells us that the universe is reionised since, at least, z » 6. Reionisation leads to a certain degree of re-scattering of CMB photons. This induces additional damping of anisotropies and additional polarisation on large scales (up to the horizon scale at reionisation). It enters the CMB spectrum mainly through one parameter, the optical depth  to the last scattering surface or the redshift of reionisation z re.

21. Gunn Peterson trough normal emission no emission In quasars with z<6.1 the photons with wavelength shorter that Ly-a are not absorbed. (from Becker et al. 2001)

22. The power spectrum of CMB anisotropies  T(n) is a function on the sphere, we can expand it in spherical harmonics observed mean cosmic variance ( if the a lm ’s are Gaussian ) consequence of statistical isotropy

23. The physics of CMB fluctuations  <  ’ 800 > Small scales : Damping of fluctuations due to the imperfect coupling of photons and electrons during recombination (Silk damping).  ’  < 1 o 100< < 800 Intermediate scales : Acoustic oscillations of the baryon/photon fluid before recombination.  > 1 o < 100 Large scales : The gravitational potential on the surface of last scattering, time dependence of the gravitational potential  ~ 10 -5.

24. Power spectra of scalar fluctuations

25. WMAP data Temperature (TT = C )Polarisation (ET) Spergel et al (2003)

26. Newer data I From Readhead et al. 2004 CBI

27. Newer data II (From T. Montroy et al. 2005) The present knowledge of the EE spectrum.

28. Observed spectrum of anisotropies Tegmark et al. 03

29. Acoustic oscillations Determine the angular distance to the last scattering surface, z 1

30. larger  Dependence on cosmological parameters more baryons Most cosmological parameters have complicated effects on the CMB spectrum

31. Geometrical degeneracy degeneracy lines : Flat Universe (ligne of constant curvature  K =0 ) Degeneracy: shift Flat Universe:  =  h 2

32. Primordial parameters Scalar spectum: scalar spectral index n S and amplitude A n S = 1 : scale invariant spectrum (Harrison-Zel’dovich) blue, n S > 1 red, n S < 1 Tensor spectum: (gravity waves) n T > 0 The ‘smoking gun’ of inflation, has not yet been detected: B modes of the polarisation (QUEST, 2006).

33. Mesured cosmological parameters Attention: FLATNESS imposed!!! On the other hand:  tot = 1.02 +/- 0.02 with the HST prior on h... z reion ~ 17 unexpectedly early reionisation (With CMB + flatness or CMB + Hubble)   =0.73 § 0.11 a rigid constraint which is in slight tension with nucleosynthesis?  bar = 0.02 + 0.002 Spergel et al. ‘03

34. Forecast1: WMAP 2 year data (Rocha et al. 2003)  b =  b h 2  m =  m h 2   =   h 2 n s spectral index Q quad. amplit. R angular diam.  optical depth

35. Forecast2: Planck 2 year data ( Rocha et al. 2003)

36. Forecast3: Cosmic variance limited data (Rocha et al. 2003)

37. Evidence for a cosmological constant Sn1a, Riess et al. 2004 (green) CMB + Hubble (orange) Bi-spectrum  Verde 2003 (blue) (from Verde, 2004)

38. The CMB is a superb, physically simple observational tool to learn more about our Universe. We know the cosmological parameters with impressive precision which will still improve considerably during the next years. We don’t understand at all the bizarre ‘mix’ of cosmic components:  b h 2 ~ 0.02,  m h 2 ~ 0.16,   ~ 0.7 The simplest model of inflation (scale invariant spectrum of scalar perturbations, vanishing curvature) is a good fit to the data. What is dark matter? What is dark energy? What is the inflaton? Conclusions ! We have not run out of problems in cosmology!