Cosmology Begin with the simplest physical system adding

Cosmology Begin with the simplest physical system, adding complexity only when required. Towards this end, Einstein introduced the Cosmological Principle : the universe is homogeneous and isotropic on sufficiently large scales… The universe looks pretty much like this everywhere – “walls” and “voids” are present but no larger structures are seen. Redshift surveys and CMB confirm this principle. It follows that the Universe has no “edge” or center – extension of Copernican Principle

Hubble’s Law (Hubble 1929) the recessional velocity of external galaxies is linearly related to their distance. recession velocity = Ho x distance Derive Hubble’s law from the Cosmological Principle Consider a small triangle. As the universe expands or contracts, the conditions of homogeneity and isotropy require that the expansion is identical in all locations. The triangle must grow self-similarly. Define the present time at to and the scale factor of the expansion as a(t), with ao = a(to) being the scale factor at to. Self-similarity requires that any distance x increase by the same scale factor. Where Hubble parameter is Ho is value of the Hubble parameter at to Note that the Cosmological Principle does not require H > 0 – can have static or contracting universe.

Dynamics of the Universe - Conservation Laws and Friedmann Equations To solve for the dynamics of the universe, we will assume the Cosmological Principle along with General Relativity. Begin with a Newtonian approximation to derive the evolution of the universe. Use Lorentzian transformations within the Newtonian framework Approximate a region of the universe as a uniform density sphere of non-relativistic matter. Use the Eulerian equations (like CBE used for stars in galaxies!) for conservation of mass and momentum to derive the dynamical evolution of the universe. + If mass is conserved, mass density satisfies EOC

Time dependence of matter density is determined solely by the evolution of the scale factor - for a matter dominated Universe ρ ~ a-3

Apply conservation of momentum using Newton’s theory of gravity (though difficult to define potential in a uniform unbound medium like the Universe!). Apply Newton’s laws to a universe which is the interior of large sphere (doesn’t violate CP badly if we consider only regions where x << Runiverse). Euler’s equation for momentum conservation or where v is the local fluid velocity, p is the pressure, and F is the force (in this case gravitational) per unit mass. If p gradient is zero (for CP) and x x = x, then, Using For F we use Poisson’s eq. And mass conservation Universe must be changing velocity if it contains matter!

Einstein used GR to improve upon Newtonian cosmology and provide a description of the Universe as a whole. He introduced the cosmological constant Λ into his GR equations since he found (as we did) that if the density of the Universe is non-zero, the Universe must be expanding. As this was before Hubble’s work and most believed in a Steady State Universe, this constant allowed for a non-zero density and a static Universe. After Hubble’s discovery of an expanding Universe, Einstein famously called this “greatest blunder”. Start with modification to Poisson’s law, Positive Λ is repulsive force that counteracts gravity. Now replace density with ρeff = ρ + 3 p/c 2 total energy density (rest mass + kinetic E) Get Friedmann solutions for the scale factor of the Universe by multiplying by da/dt and integrating wrt time. For expanding Universe, gravitational term dominates when scale factor is small. At later times, first the curvature term then later the cosmological constant dominate.

Let’s work through one possible solution to the Friedmann equations with a simplicity assumption of a null cosmological constant, Λ = 0, and recall Evaluate the constant k in terms of present day observable quantities. k = 0 only when density is at the critical density, defined as Cosmological Density Parameter Inserting back into Friedmann equation we get

Case 1: k = 0 and Ωo= 1 Einstein-de. Sitter universe Case 2: k > 0 and Ωo> 1 Case 3: k < 0 and Ωo< 1 As scale factor increases, it eventually reaches a point where adot = 0. Expansion stops at amax. Energy term is now positive. Solution for a(t) is analogous to a rocket shot with velocity greater than escape velocity. a(t) is proportional to t The Universe expands at an ever decreasing rate! After amax is reached, Universe starts to collapse! Borderline Universe or Marginally Bound Closed or Bound Universe Expansion continues forever! Open or Unbound Universe

The Fate of the Universe If Λ = 0, the fate of the universe (whether it continues to expand forever or eventually collapses) depends on its density. High density = lots of mass = enough matter to gravitationally halt expansion and cause gravitational collapse Low density = little mass = not enough gravitational attraction to stop expansion…it goes on forever If Ho = 70 km/s/Mpc, then ρcrit ≈ 1 x 10 -29 g/cm 3 (about one H atom in 200 L volume of space)

Cosmology and General Relativity Einstein’s Principle of Equivalence space-time is curved in the presence of a gravitational field (the mass density of the Universe tells us about the geometry of space-time) Interpret the solutions for the scale factor obtained from Newtonian theory within the framework of GR (Friedmann equations can be derived from Einstein’s field equations). k is positive or o > 1, Bound/Closed Universe k is negative or o < 1, Open Universe k = 0 or o = 1, Flat/Marginally Bound/Critical Universe

Use 2 D analogy to visualize curvature of space-time • Detected by measuring the sum of the angles of any triangle • Locally, space may appear flat (Euclidean) • Define a metric along the surface A metric, or distance measure, describes the distance, ds, between two points in space or space-time. For a unit sphere A geodesic is the shortest distance between two points and can be found by minimizing the space or space-time interval, Space-time metric: For special relativity, in a Lorentz frame we define a distance in spacetime as For light, the metric yields ds 2 = 0. Light follows a null geodesic - meaning that the physical distance travelled is equal to ct. Everything that we see in the universe by definition lies along null geodesics, as the light has just had enough time to reach us.

3 Types of spacetime intervals: ds 2 = zero - lightlike (photons travel along these lines - this is a lightcone in 2 -d space) ds 2 > zero - timelike (positions are close enough in space that a photon would have had more than enough time to travel from one event to the other) ds 2 < zero - spacelike (photon cannot traverse the distance in the time given - one event could not have caused the other)

To construct a metric that is valid in a cosmological context, assume • the cosmological principle is true • each point in spacetime has one and only one co-moving, timelike geodesic passing through it If the expansion of the universe is homogeneous and isotropic, co-moving coordinates are constant with time. For a co-moving observer, there is a metric for the universe called the Robertson-Walker metric, or the Friedmann-Le. Maitre-Robertson-Walker metric Where a is the scale factor of the Universe, r, θ, ϕ are the co-moving coordinates, k is the sign of curvature (0, -1, +1) and dη is the solid angle. This metric specifies the geometry of the universe with just one undetermined factor, a(t), which is determined from the Friedmann equations. Together, the Friedmann equations and RW metric completely describe the geometry. Co-moving distance DC = ro Proper distance DP = (a/ao)ro (Proper distance does change with time)

Cosmological Redshift We now have a way to describe the evolution of the size of the universe (Friedmann equation) and measuring distances within the universe (Robertson. Walker metric). Recast these in terms of observable quantities we don’t directly observe a(t), but we can observe the cosmological redshift of objects due to the expansion of the universe. Recall the Doppler shift of light (redshift or blueshift) is defined as Since ao is a(to) then, ao=1 As Universe expands, photon’s λ expands proportionally to the scale factor a(t) Since redshift arises due to expanding wavelength of all photons traveling through an expanding Universe, it is called a cosmological redshift. As a consequence, we don’t normally convert redshift to distance, since we need to assume a particular model for how a(t) has evolved.

Friedmann equation in observable quantities Friedmann equation Critical value of Λ in flat, empty universe Rewrite Friedmann with these substitutions At time to Fundamental component of the Friedmann equation upon which our measures for the distance and evolution of other quantities will be based.

Expressing Distances in an Expanding Universe The geometry and expansion rate of the Universe effects angular sizes and distances measured. Integrate over components of RW metric. DH = c/Ho Hubble Distance (distance light travels in Hubble time, t. H = 1/Ho) DC= Radial Co-moving Distance DM = DC (flat) Transverse Co-moving Distance, differs for curved space (see Hogg 2000) DA = L(proper length)/θ(angular size) = DM/(1+z) Angular Distance DL = sqrt (L/4π*flux) = DM(1+z) = DA(1+z)2 Luminosity Distance If Λ = 0 and flat geometry, then DL = 2 c/Ho [z/(G+1)] {1+[z/(G+1)]} where G = (1 + z)1/2 See Ned Wright’s Javascript Cosmology Calculator for DL for different cosmologies: http: //www. astro. ucla. edu/~wright/Cosmo. Calc. html

Angular Size vs Redshift As an object is moved to higher redshift, its angular size first decreases (as expected) but soon begins to increase after passing through a minimum value. The appearance of a minimum angular size at a given redshift zmin is a generic feature of cosmological models with Ωm > 0. Figure 10 from Sahni and Starobinsky (2000) The angular size as a function of cosmological redshift z for flat cosmological models ΩM+ΩΛ= 1. Heavier lines correspond to larger values of ΩM.

Angular diameter distance vs z (plotting DA/DH where DA=L/θ) Luminosity distance vs z (plotting DL/DH) DH=c/Ho= 3000 h-1 Mpc At high z, angular diameter distance is such that 1 arcsec is about 5 kpc. flat, Λ=0 – solid open, Λ=0 – dotted flat, non-zero Λ - dashed (from Hogg 2000 astro-ph 9905116)

Surface Brightness Dimming SB is flux per unit solid angle Recall Flux = L/(4πDL 2) and the solid angle subtended by a source of projected area A is Ω = A/DA 2 Since DL = DA(1+z)2 we can show the surface brightness is Observed surface brightness of objects decreases very rapidly as one moves to high redshifts purely due to cosmology. Cosmological dimming is independent of cosmological parameters.

How does the age of the Universe differ in different cosmologies? Friedman models (Λ = 0) give to = (2/3)Ho-1 0 < to < (2/3)Ho-1 < to < Ho-1 k=0 k = +1 k = -1 What is the relationship between redshift and age of the Universe? Lookback time (how long ago photons left object at z) t. L = 2/3 Ho-1[1 -(1+z)-3/2] for flat Universe Age of the Universe when light left source at redshift z te(Gyr) = 10. 5 (Ho/65) (1+z)-3/2 for flat Universe

Lookback times vs z plotting t. L/t. H and age vs z plotting t/t. H t. L is difference between age of Universe now and age t. U when photons left emitting source flat, Λ=0 – solid open, Λ=0 – dotted flat, non-zero Λ - dashed

Age of the Universe as calculated from various models. A flat universe with a significant cosmological constant (ΩM=0. 3, ΩΛ=0. 7) yields an age close to what you get with a constant value of Ho (t. H or 1/Ho). 14 Age of ΩM=1, ΩΛ=0 universe is about 9 billion years Age of ΩM=0. 3, ΩΛ=0. 7 universe is about 14 billion years

Comparison of Different Cosmological Models qo is “deceleration parameter”

Classical Cosmological Tests (or How to determine if the Universe is Open, Closed, Flat or Accelerating? ) 1) Add up all matter in the Universe to determine mass density Luminous matter: only ~1% of ρcrit Dark matter? Still a factor of ~5 to low to close the Universe with DM This only constrains ΩM and not the value of the cosmological constant ΩΛ which also impacts fate/age of the Universe From Big Bang Nucleosynthesis, we will see that baryonic matter is about 45% of the critical density with all matter (baryonic+DM) totaling ~30%.

2) Measure the curvature of space-time by surveying the Universe on large scales. Volume of space is a function of cosmological parameters. Thus, for a given class of objects the redshift distribution, N(z), will depend upon Ω 0 (or ΩM) and ΩΛ. Assume a uniformly distributed population of objects with mean density n 0. Then, Measurements attempted with QSOs, though must account for evolutionary effects. Results indicate that if the Universe is curved, radius of curvature would have to be comparable to or greater than the Hubble distance – essentially flat…but why…. This is a challenge for Big Bang Cosmology

3) Look for dynamical effects on the expansion rate of the Universe (impacted by both ΩM and ΩΛ ). Models showing unbound, decelerating universe – more distant objects should be moving faster than nearby objects if the universe is decelerating. Objects were receding less rapidly in the past! 70 km/s/Mpc The rate of cosmic expansion can be determined by probing objects at great distances where the geometry and Hubble parameter changes can be detected. Use Type 1 supernovae – standard candles that can be seen in distant galaxies (several Gpc) galaxy distance can be determined independent from Universal expansion.

Results from Type 1 a Supernovae surveys: First evidence for accelerating Universal expansion Flatness criterion – dashed line Acceleration =0 - solid diagonal Big Crunch below solid horizontal