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Cosmic Microwave Background (CMB) Cosmic Microwavepenzias and Wilson Background (1965) an orientation and an invitation we are here photo Courtesy: A Frolov, Aspen, CO 2017 Image Credit: Hubble Ultra-Deep Field, NASA/ESA we are here Length Scales ? few kpc ⇠ few light hours ⇠ few m ⇠ 9 few thousand km Gpc = 10 pc Mpc = 106 pc kpc = 103 pc pc 3.26 lyrs ⇡ few Mpc ⇠ Image Credit: Hubble Ultra-Deep Field, NASA/ESA large scale distribution of galaxies few Gpc ⇠ Image Credit: Hubble Ultra-Deep Field, NASA/ESA large scale isotropy and homogeneity (+ inhomogeneities) few Gpc ⇠ cosmological principle: no special location or directionImage Credit: Hubble Ultra-Deep Field, NASA/ESA dynamics: what are distant galaxies doing ? Image Credit: Hubble Ultra-Deep Field, NASA/ESA 1930’s interpreted as the universe is expanding separation a(t2)∆x separation a(t1)∆x time a(t)=scalefactor our cosmic story 1 1 a˙(today) = H 68 km s− Mpc− 0 ⇡ g expandin 1 10 age of universe H− 10 yrs ⇠ 0 ⇠ credit: adapted from NASA early universe was hotter expanding (filled with radiationhot + matter) cold credit: adapted from NASA too hot for atoms to exist - - + - + + + - + - - - + - + - + - - + + + - + - Cosmic Microwave Background (CMB) cosmic microwavePenzias and Wilson background (1965) T0 =2.73K Credit: Roger Ressmeyer cosmic microwave background T =2.726 0.001 K ± COBE/FIRAS Mather et. al 1994 Fixen et. al 1996 Cosmic Microwave Background temperature anisotropies 5 δT/T 10− ⇠ Planck 2015 today the deviations from homogeneity are larger! Image Credit: Hubble Ultra-Deep Field, NASA/ESA our cosmic story lumpier homogeneous g expandin hot (filled with radiation + matter) cold credit: adapted from NASA growth of inhomogeneities ? formation of structure expanding there is more going on: chemical enrichment primordial late our complete cosmic story? formation of structure universe gets! populated first atoms with particles first nuclei Reheating ! protons & neutrons after inflation Cosmology@ Inflation ? ? ?? << sec few min 380,000 yrs 14 billion years seems sufficient ? + Quantum Field Theory + General+ Relativity not so fast … galaxies are rotating too fast … requires additional unseen matter assuming Newton’s/Einstein’s gravity galaxies cluster are way too hot … requires additional unseen matter hot cluster: x-rays! too much distortion of images … requires additional unseen matter CMB & growth of structure … requires additional matter that does not interact with much with light so we need some “dark matter” … and then there is … universe is undergoing accelerated expansion ! distances from Type Ia Supernovae March 14, 1997 May 18, 1998 universe is undergoing accelerated expansion ! a¨(t) (⇢ +3P ) / require something with significant negative pressure! (Dark Energy) our universe ~68 % DARK ENERGY ~27% DARKMATTER STARS, GAS, ETC 5% and … why is this so uniform? requires seemingly causal correlations T =2.726 0.001 K ± COBE/FIRAS Mather et. al 1994 Fixen et. al 1996 Cosmic Microwave Background origin of temperature anisotropies ? 5 δT/T 10− ⇠ Planck 2015 quantum fluctuation + accelerated expansion (inflation) = origin of anisotropies separation a>¨ 0 a eHt ⇠ a>¨ 0 quantum jitter in space-time but there is more … you will learn: - how we know what we know about the cosmos - our best physical model for our universe along with “known” physics, we will need to add - dark matter - dark energy - inflation amongst other things, also matter-antimatter asymmetry 125 6. Initial Conditions from Inflation non-linear 124 6. Initial Conditions from Inflation Galaxy Clustering (Reid et al. 2010) CMB (Hlozek et al. 2011) Exercise.—Show that LyA (McDonald et al. 2006) r = 16" (6.5.77) CMB Lensing (Das et al. 2011) linear (reconstructed) n = 2" . (6.5.78) Clusters (Sehgal et al. 2011) t − Weak Lensing (Tinker et al. 2011) precision cosmologyNotice ! that this implies the consistency relation nt = r/8. − Inflationary models can be classified according to their predictions for the parameters ns and r. Fig. 6.3 shows the predictions of various slow-roll models as well as the latest constraints Figure 6.4: Compilation of the latest measurements of the matter power spectrum.from measurements of the Planck satellite. 0.25 6000 0.20 concave 5000 convex 0.15 4000 small-field 0.10 large-field 3000 2000 0.05 natural 1000 26 1. Geometry and Dynamics 0.00 0.940 94 0.960 96 0.980 98 1.00 500 200 74 3. Thermal HistoryFigure 6.3: Latest constraints on the scalar spectral index ns and the tensor amplitude r. 26 HST 0.32 0.68 250 100 1.00 0.00 24 SNLS 0 0 -250 -100 22 SDSS 6.6 Observations -500 -200 WMAP 20 2 5 10 20 500 1000 1500 2000 Inflation2500 predicts nearly scale-invariant spectra of superhorizon scalar and tensor fluctuations. 18 Once these modes enter the horizon, they start to evolve according to the processes described in 125 6. Initial Conditions from Inflation Chapter 5. Since we understand the physics of the subhorizon evolution very well, we can use 16 Low-z Figure 6.5: The latest measurements of the CMB angular power spectrum by the Plancklate-time satellite. observations to learn about the initial conditions. distance (apparent magnitude) distance (apparent 14 4 0.00.2 0.4 0.6 0.8 1.0 1.2 1.46.6.2 CMB Anisotropies 6.6.1 Matter PowerMass Fraction Spectrum He redshift In Chapter 5, we showed that subhorizon perturbations evolve di↵erently in the radiation- The temperature fluctuations in the cosmic microwave background are sourced predominantly Figure 1.9: Type IA supernovae and the discovery dark energy. If we assume a flat universe, then the dominated and matter-dominated epochs. We have seen how this leads to a characteristic shape supernovae clearly appear fainter (or more distant) than predicted in a matter-onlyby universe scalar (⌦m (density)=1.0). fluctuations. Acoustic oscillations in the primordial plasmaof the before matter recombi- power spectrum, cf. fig. 5.4. In fig. 6.4 we compare this prediction to the measured (SDSS = Sloan Digital Sky Survey; SNLS = SuperNova Legacy Survey; HST = Hubble Space Telescope.) D nation lead to a characteristic peak structure of the angular power spectrummatter of the power. CMB;5 see 0.80 fig. 6.5. The precise shape of the spectrum depends both on thenon-linear initial conditions5 With (through the exception the of gravitational lensing, we unfortunately never observe the dark matter directly. Instead +lensing 75 parameters A and n ) and the cosmological parameters (through parametersgalaxy like surveys⌦ , like⌦ the, ⌦ Sloan, Digital Sky Survey (SDSS)3 only probe luminous matter. On large scales, the density +lensing+BAO s s m ⇤ k He 70 contrast for galaxies, ∆g, is simply proportional to density contrast for dark matter: ∆g = b∆m, where the bias 0.72 parameter b is a constant. On small scales, the relationship isn’t as simple. 65 Galaxy Clustering (Reid et al. 2010) 60 CMB (Hlozek et al. 2011) 0.64 55 LyA (McDonald et al. 2006) CMB Lensing (Das et al. 2011) linear 7 50 (reconstructed) to H Number relative Li 0.56 Clusters (Sehgal et al. 2011) 45 Weak Lensing (Tinker et al. 2011) 40 0.24 0.32 0.40 0.48 Figure 1.10: AcombinationCMBandLSSobservationsindicatethatthespatialgeometryoftheuniverse WMAP is flat. The energy density of the universe is dominated by a cosmological constant. Notice that the CMB Figure 6.4: Compilation of the latest measurements of the matter power spectrum. data alone cannot exclude a matter-only universe with large spatial curvature. The evidence for dark energy Figure 3.10: Theoretical predictions (colored bands) and observational constraints (grey bands). requires additional input. ⌧ : a large neutron lifetime would reduce the amount of neutron decay after freeze-out Single-Component Universe 6000 • n 4 3 0 and therefore would increase the final helium abundance. The di↵erent scalings of radiation (a− ), matter (a− ) and vacuum energy (a ) imply that for 5000 most of its history the universe was dominated by a single component (first radiation, then : a larger mass di↵erence between neutrons and protons would decrease the n/p ratio matter, then vacuum energy; see fig. 1.11). Parameterising this component by its equation4000 of • Q state wI captures all cases of interest. For a flat, single-component universe, the Friedmann at freeze-out and therefore would decrease the final helium abundance. equation (1.3.135)reducesto 3000 a˙ 3 (1+wI ) ⌘: the amount of helium increases with increasing ⌘ as nucleosythesis starts earlier for = H0 ⌦I a− 2 . (1.3.136) • a 2000 p larger baryon density. 1000 GN : increasing the strength of gravity would increase the freeze-out temperature, Tf • 1/6 / GN , and hence would increase the final helium abundance. 500 200 2/3 250 G : increasing100 the weak force would decrease the freeze-out temperature, T G− , • F f / F 0 and hence would0 decrease the final helium abundance. -250 -100 Changing the input, e.g. by new physics beyond the Standard Model (BSM) in the early universe, -500 -200 would change the predictions of BBN. In this way BBN is a probe of fundamental physics. 2 5 10 20 500 1000 1500 2000 2500 Light Element Synthesis⇤ Figure 6.5: The latest measurements of the CMB angular power spectrum by the Planck satellite. To determine the abundances of other light elements, the coupled Boltzmann equations have to 6.6.2 CMB Anisotropies be solved numerically (see fig. 3.11 for the result of such a computation). Fig. 3.10 shows that theoretical predictions for the light element abundances as a function of ⌘ (or ⌦b). The fact that The temperature fluctuations in the cosmic microwave backgroundwe find are sourced reasonably predominantly good quantitative agreement with observations is one of the great triumphs by scalar (density) fluctuations. Acoustic oscillations in the primordialof the Big plasma Bang before model.
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