Big Bang Nucleosynthesis Finally, relative abundances are sensitive to density of normal (baryonic matter)

Thus Ωb,0 ~ 4%. So our universe

Ωtotal ~1 with 70% in Dark Energy, 30% in matter but only 4% baryonic! Case for the Hot Big Bang

• The Cosmic Microwave Background has an isotropic blackbody spectrum – it is extremely difficult to generate a blackbody background in other models

• The observed abundances of the light isotopes are reasonably consistent with predictions – again, a hot initial state is the natural way to generate these

• Many astrophysical populations (e.g. quasars) show strong evolution with redshift – this certainly argues against any Steady State models The Accelerating Universe

Distant SNe appear too faint, must be further away than in a non-accelerating universe.

Perlmutter et al. 2003 Riese 2000 Outstanding problems

• Why is the CMB so isotropic? – horizon distance at last scattering << horizon distance now – why would causally disconnected regions have the same temperature to 1 part in 105?

• Why is universe so flat? – if Ω is not 1, Ω evolves rapidly away from 1 in radiation or matter dominated universe – but CMB analysis shows Ω = 1 to high accuracy – so either Ω=1 (why?) or Ω is fine tuned to very nearly 1

• How do structures form? – if early universe is so very nearly uniform Astronomy 422

Lecture 22: Early Universe Key concepts:

Problems with Hot Big Bang Inflation

Announcements:

April 26: Exam 3 April 28: Presentations begin Astro 422 Presentations:

Thursday April 28: 9:30 – 9:50 _Isaiah Santistevan______9:50 – 10:10 _Cameron Trapp______10:10 – 10:30 _Jessica Lopez______Tuesday May 3: 9:30 – 9:50 __Chris Quintana______9:50 – 10:10 __Austin Vaitkus______10:10 – 10:30 __Kathryn Jackson______Thursday May 5: 9:30 – 9:50 _Montie Avery______9:50 – 10:10 _Andrea Tallbrother______10:10 – 10:30 _Veronica Dike______10:30 – 10:50 _Kirtus Leyba______

Send me your preference. First come, first served. Main goals

• With cosmology we aim to get a physical description of the universe. – What is the matter content? – What is the dynamics of the universe?

• We have worked on a set of equations describing the universe, and we introduced a number of cosmological parameters. – Hubble parameter, scale factor, density parameter etc.

• We would like to measure these parameters to get the theories right.

• Eventually, we use the models to understand the origin and evolution of cosmic structure.

Epochs

• Friedmann's equation can be written as (Eq. 29.122):

• At very early times, when R<<1, only the first two terms are important.

• Radiation dominates until z~40000

• CMB released at z~1100 which is well into the matter dominated era. Planck units

• GR expected to break down as Big Bang is approached

• The estimate for onset of quantum mechanical effects on spacetime is when the de Broglie wavelength equals the Schwarzschild radius:

• This defines what we call the Planck mass • We then define the Planck Length :

• and the corresponding Planck time:

• This corresponds to the earliest time that we can address by current physical theory. Our theories include a set of fundamental constants, and these are the limits of current theories. Horizon problem Start with the Horizon distance (recap)

• Universe expands and ages, an observer at any point is able to see more distant objects (light has time to arrive)

• As time increases, larger and larger regions of the universe come into causal contact with the observer

• The proper distance to the farthest observable point at a time t is called

the horizon distance, dh.

– Two points separated by a distance > dh are not in causal contact

– dh is the diameter of the largest causally connected region. Illustration: causal contact

Insufficient time for signal to reach central observer

Consider 3 locations in space, A, B & C. A&B are causally connected.

A B C Particle horizon

Causality problem: how can structures on large scale in the CMB know about each other? How did the temperature get so uniform? Flatness

Friedmann's equation again

implying

during the matter-dominated era, we had • Then we achieve

i.e., deviation of the total density parameter from unity grows with time.

• Fine-tuning problem: example, if today

At the Planck time,

This is what we refer to as the flatness problem. Structure problem

• Structure in the universe (galaxies, clusters of galaxies etc) came from inhomogeneities in the early universe.

• We see those same inhomogeneities in the CMB maps

• How was this coherence achieved? • How did the structures get there?

• Why are they just the right magnitude and size to produce the structures we see today?

• How is it possible to have the same kind of inhomogeneities spread through the whole universe without causal contact between different parts of the early universe?

– CMB is statistically the same in all directions

– Galaxies that are formed are similar in properties on opposite sides of the universe The relic problem

• Analogy: consider cooling of a liquid like water

• Once liquid reaches freezing point: – freezing does not occur smoothly and uniformly – starts at certain locations, and the crystals starts growing – when crystals merge to form solid, there will be dislocations where individual crystals meet

• The process of freezing is called a phase transition, that is when matter is changing from one phase to another. Example: dislocations in steel • This could produce exotic structures that we call topological effects – domain walls (2d sheet-like structures) – cosmic strings (1d string-like structures) – None of these have been seen in the observable universe (good limits from CMB data: strings would gravitationally lens the background)

• GUTs also predict exotic particles produced in the early universe – magnetic monopoles – never detected, and don't reveal their presence in any observed phenomena

• The absence of monopoles (and other relics predicted by particle physics theories) is called the relic problem. Summary of problems with the Hot Big Bang model: • Horizon problem – how could the CMB acquire a single temperature across the sky? – causality: how can structures in the CMB know about each other?

• Flatness problem – how could the density parameter be fine-tuned such that today?

• The relic problem – where are all the magnetic monopoles predicted?

• Origin of structure – How could structure arise in this very smooth universe

In principle, the purpose of inflation is to create a large, flat, homogenous universe. • Theory of cosmic inflation was first suggested by Alan Guth in 1982

• He postulated an inflationary epoch – very rapid, exponential expansion of the universe – occurs during the first 10-37-10-32 sec – During this time, the universe expanded by a factor of 1040-10100

• Inflation and the radius of the observable universe That would correspond to a drop in the density parameter

time

Start of End of Now Distant inflation inflation future Timeline:

Does this rapid expansion imply a violation of relativity?

No, because it is space itself that is expanding, R(t), rather than material particles moving apart at high speed in a fixed, stationary space.

Solving cosmological problems with inflation

• The flatness problem: – take any reasonably curved surface – expand it by enormous factor – after expansion it look flat locally – so inflation predicts a universe that is indistinguishable from being flat The horizon problem (without inflation) How does inflation solve the horizon problem?

• Before inflation, the particle horizon is about 10-29m – determines causally connected region

• After inflation, the particle horizon is about 1011 (up to 1070) m – normal expansion takes over, and expands by another factor 1022 before the epoch of last scattering surface – Thus causally connected regions have sizes of at least 1033m!

• Since decoupling, the scale factor have increased by at least a factor of 1000 (since z=1100 at decoupling), so the causally connected radius would be 1036 m – Current horizon of the universe is about 1026 m. – Thus, the observable universe is a small part of the causally connected part of the early universe!

How inflation solves the structure problem:

• Initial inhomogeneities due to quantum fluctuations during inflationary epoch.

• Virtual particle pairs that formed would be separated by inflationary expansion before they could annihilate, creating uneven densities. • Inhomogeneities were continually created, and then stretched to much larger scales, outside the horizon

• This naturally gives a characteristic power spectrum for the inhomogeneities

• Fluctuations created by inflation can only grow at much later times – the particle horizon must expand so that is larger than the size scale of that fluctuation – since the particle horizon increases with time, the smaller scale fluctuations grow first – clusters, voids, are results of quantum fluctuations originally occurring on tiny scales! How inflation solves the relic problem

• Suppose you have exotic particles, or structures (cosmic strings, magnetic monopoles etc) created in the very early universe

• They would become very diluted during the inflationary epoch, since the space expanded enormously – The probability that we see a relic exotic particle is then very small.

• Hm, what about baryons then? Wouldn’t there be a very small probability of finding them as well? – No, because baryogenesis occurred after inflation, it was too hot before. Vacuum energy is converted to regular matter, including baryons, and radiation. Why did inflation happen?

The answer is thought to be in behavior of quantum fields - we're in a bubble universe.

Lets look at some quantum description of radiation: • EM field is the basic entity, and it permeates all space – Photons are considered excitations, or ripples, of the field with certain wavelengths and frequencies – Energy and momentum of the excitations in the field are quantized

• Every particle has its 'own' field – electron field, quark field, gluon field (corresponding excitations are electrons, quarks, and gluons)

• Recall Heisenberg: position and momentum cannot be known simultaneously, but obey some probabilistic rules related to field The inflation model (or the basics of most of them) • Original idea by Alan Guth – Some particle called inflaton existed with a corresponding quantum field – inflation occurs during transition from a false to a true vacuum, releasing a 'false vacuum' energy that drove inflation (cf dark energy driving expansion now) – eventually field is unstuck and evolves to a lower energy state, which is the 'true vacuum', and then inflation ends. • During inflation, temperature drops because T is inversely proportional to R

• After inflation, vacuum energy is converted to matter and radiation – this heats the universe, so T increases

• The following evolution is the radiation dominated, followed by the matter dominated etc.

BICEP2 results

Measurement of B-Modes by BICEP2 If primordial these could constrain inflation models.

Planck should be able to confirm these results (sorry, its just dust).

B-Modes

BICEP2 in the Antarctic Chaotic inflation

• Idea is that inflation occurs due to fluctuations in some kind of quantum field in the early universe

• Some regions inflate and some don't – our whole observable universe is a sub-part of one of the bubbles that did inflate

• A larger universe (super universe) may be continually spawning new bubble universes within it – may be other interesting bubbles out there, but we will never be able to observe or communicate with them since they are outside our horizon – Unless we get hit by an inflating bubble?