# LOG#112. Basic Cosmology (VII).

**Posted:**2013/06/24

**Filed under:**Cosmology, Physmatics |

**Tags:**cosmological parameters, inflaton field, PLANCK results, power spectrum, scalar perturbations, spectral index, tensor perturbations, WMAP results Leave a comment

In my final post of this basic Cosmology thread, I am going to discuss scalar perturbations and to review some of the recent results about cosmological parameters by WMAP and PLANCK.

One of the predictions of the Standard Cosmological Model (LCDM or **Λ**CDM) is that we should expect some inhomogeneities and anisotropies in the Cosmic Microwave Background (CMB). Indeed, the “perturbations” of densities (matter, radiation,…) are closely related to these anisotropies and the easiest way to trigger them is by using “scalar perturbations”, i.e., perturbations induced by the scalar field in the early Universe! In fact, inflation should be understood at quantum level as a “quantum fluctuation” that triggers scalar perturbations of the so called** “inflaton field”.**

The scalar field perturbations are “shifts” or “fluctuations” in the scalar field, namely

Now, we have to work out a little bit the equation for , if we take into account that our Universe is seemingly a “smoothly” expanding Universe. To first order, we can neglect the “back reaction” on the metric field due to the scalar field . Therefore, the equation we get is really simple:

or equivalently

where

since the term can be neglected during inflation. Furthermore, we can write

so

for any Fourier mode “k” in the Fourier decomposition of the scalar field . If we define , then we can rewrite the above equation as follows

This last equation describes** a harmonic oscillator equation with time-varying frequency**! We can quantize using the formalism of creation and annihilation operators (, respectively):

and where U satisfies the equation of a harmonic oscillator as well

The variance of perturbations in are defined with the equations

and where we have defined the power spectrum in terms of correlation function as the quantity

We require a solution for the equation above for . During inflation, provides

and

Then, for the equation above for U becomes

And now, a subtle issue…Initial conditions! Any initial condition (with proper normalization) long before the horizon exit implies that . It implies that the term dominates and we can substitute the last equation by this one

and this equation has pretty simple solutions with proper normalization. Essentially, they are “plane waves” modulated with the Fourier mode k, i.e.,

A more precise proper solution is

Please, note that if , we recover the previous equation. After many e-folds, we get and

Thus, the power spectrum is deduced to be

For tensor perturbations (I have not discussed them here for simplicity), this result can be compared with

So, at least, scalar perturbations have the same order of magnitude and the same power of the Fourier modes! Indeed, the fluctuations are transfered to scalar perturbations or metric pertubations, since

After the horizon crossing, we obtain

Now, we can introduce the important concept of** spectral index.** For scalar fields, the perturbations are defined in terms of this spectral index as follows

For metric field perturbations, we write

and where the spectral index and its tensor analogue are defined through the simple expressions

## Final conclusions, WMAP and PLANCK

The history of the Universe up to the period of BBN is (more or less) well understood in the framework of the Standard Cosmological Model (LCDM). It is well established and tested, BUT, please, do not forget we have to choose “the proper initial conditions” in order to obtain the right (measured/observed) results. In summary, up to the BBN we get

1. **A homogenous and isotropic Universe, with small density perturbations.** It is equivalent to a thermal bath with temperature above 1MeV, i.e., . Indeed, inflation seems to be the simplest and best solution to many problems of modern cosmology (e.g., the horizon problem or the flatness problem). However, the LCDM and the inflation itself can not answer what the inflaton field is and why the anisotropies are so tiny as we observed today. Moreover, the cosmological constant problem remains unanswered today (even when we have just discovered the Higgs field, and that some people are speculating that the Higgs field could be matched with the inflaton field, this idea has some technical problems).

2. **The energy-matter content of the Universe**…There is a wide agreement about the current composition of the Universe, based on the “concordance model” by several experiments, now confirmed by the WMAP/PLANCK probes. Generally speaking:

–** Baryons** are negligible since . In fact, only about a 5% is “matter” we do know through the Standard Model of particle physics. Most of that 5% is “light” (a.k.a., photons) and the remaining 1% are the elements we do know from the Periodic Table (and mainly hydrogen and helium, of course).

–**(Cold) Dark Matter** is about . There is no known particle that can be assigned to this dark matter stuff. It does not emit light, it is (likely) non-relativistic and uncharged under the gauge symmetry group of the SM, i.e., under . Of course, you could avoid the need for this stuff invoking a modified newtonian dynamics (MOND) and/or modified gravity (MOG). However, to my knowledge, any MOND/MOG theory generally ends by requiring some kind of stuff that is “like cold DM”. So, Occam’s razor here punches and says that cold DM is (likely) the most promising solution to the flat rotation curve of galaxies and this mysterious energy density. Remember: Dark Matter can not be any known particle belonging to the Standard Model we know today (circa 2013).

–**Dark Energy (DE)**, a.k.a. vacuum energy or the cosmological constant. Reintroduced in 1998 to explain the SNae results and the high redshifts we measured from them, its existence is now fully established as well. Perhaps the dark energy name is “unfortunated” since it is, indeed, a negative pressure terms in the Einstein’s field equations. It is about the 70 or 80% of the current energy-matter budget of the Universe, i.e., it has an energy density about .

Let me draw again a picture I have showed before in this blog:

The main results for the cosmological parameters can be found here (in the case of the WMAP9 mission): http://arxiv.org/pdf/1212.5226v3.pdf

The main results for the cosmological parameter can be found here (in the case of the Planck mission): http://arxiv.org/pdf/1303.5076v1.pdf

We can compare the cosmological parameters coming from the WMAP9 mission

with those of the Planck mission

We observe that the different experiments are providing convergent values of the main cosmological parameters, so we know we are in the right track!

Cosmology is a fascinating subject. I have only sketched some elementary ideas here. I will be discussing more advanced topics in the near future, but you have to wait for it! :).

See you in my next blog post!

# LOG#111. Basic Cosmology (VI).

**Posted:**2013/06/23

**Filed under:**Cosmology, Physmatics |

**Tags:**10-folds, cosmological constant problem, e-folds, flatness problem, horizon problem, inflation, Issues of Big Bang Cosmology, negative pressure, phantom energy, quintessence, scalar field Leave a comment

The topic today: problems in the Standard Cosmological Model (LCDM), inflation and scalar fields!

## STANDARD COSMOLOGICAL MODEL: ISSUES

Despite the success of the Standard Cosmological model (or LCDM), today it is widely accepted that it is not complete, even if its main features and observables are known, there are some questions we can not understand in that framework.

Firstly, we have the **horizon problem. **For any comoving horizon, we obtain

Note that for a MD Universe and to for a RD Universe.

According to the CMB observations, today our Universe is very close to be isotropic, since the deviations are tiny

The issue, of course, is how can it be true?Recall that the largest scales observed today have entered theo horizon just recently, long after decoupling. Moreover, microscopic causal physics can not make it! So, where the above tiny anisotropy comes from? E.g., for a distant galaxy, we get

and this scale was outside the horizon in the past! Please, note that the entropy within a horizon volume is about

and so

Another problem is the **flatness problem. **The physical radius of curvature is given by the term

The total energy density as a function of the scale factor is given by

When the primordial nucleosynthesis happened, about , it gave and it gives . By the other hand, at the Planck time, i.e., if , it gives the value , and then . This large mismatch means that the Big Bang Universe requires VERY SPECIAL initial conditions, otherwise the Universe would not be (apparently) flat as we observe at current time. Note that if and at Planck time, then there are two options:

1) k>0: the Universe recollapse withim few .

2) k<0: the Universe reaches 3K at .

Therefore, the natural time scale for cosmology is the Planck time (at least in the very early Universe and before) . The current age of the Universe is about .

An important, yet unsolved and terrific, problem is the **cosmological constant problem. **The vacuum energy we observe today (in the form of dark energy) is given by

The equation of state for the vacuum energy is known to be , i.e., . It yields

i.e., the so called de Sitter space (dS) or de Sitter Universe. It is a maximally symmetric spacetime. In a dS Universe, we obtain

Observations provide (via dark energy) that

Theoretical (and “natural”) Quantum Field Theory (QFT) calculations give

or so( even the mismatch can be 122 or 123 orders of magnitude!!!!). This problem is far beyond our current knowledge of QFT. It (likely) requires new physics or to rethink QFT and/or the observed value of the vacuum energy density. It is a hint that our understanding of the Universe is not complete.

## INFLATION

One of the most “simple” and elegant solutions to the flatness problem and the horizon problem is the inflationary theory. What is inflation? Let me explain it here better. If the (early) Universe experienced an stage of “very fast” expansion, we can solve the horizon problem! However, there is a problem (you can call Houston if you want to…). If we do want a quick expansion in the early Universe, we need “negative pressure” to realize that scenario. Negative pressure can be obtained in an scalar field theory (there are some alternatives with a repulsive vector field and/or higher order tensors like a 3-form antisymmetric field, but the most simple solution is given by scalar fields).

The solution to the horizon problem in the inflationary theory proceeds as follows. Firstly, for the comoving horizon, we get

where the disance over which particles can travel in the course of one expansion time is the comoving Hubble radius , and it is encoded into the fraction

We have to make a distinction of the comoving horizon and the comoving Hubble radius . If we observe two particles whose comoving distance is r, then

1st. If , we can never have a communication between those particles.

2nd. If , then we can never communicate NOW using those two particles.

It shows that it is possible to have . A particle with ca not communicate today BUT they could be in causal contact early on. We only need that . That is, get contributions mostly from early epochs. Indeed, both in RD (Radiation Dominated) and MD(Matter Dominated) Universes, increase with time, so the latter epoch contributions dominate over cosmological time scales. Then, a solution for the horizon problem is that in the early Universe, in this “inflationary” (very fast) phase, for at least a brief amount of time (how much is not clear) the comoving Hubble radius DECREASED.

How can the scale factor evolve in order to solve the horizon problem?We do know that must decrease, so must increase! Therefore,

i.e., we get an (positively) accelerating expansion or **“inflation”. **

How can we understand quantitatively inflation? Firstly, suppose that the energy scale of inflation is about . It is about the Grand Unified Theory (GUT) energy scale, and close to the Quantum Gravity scale (the Planck scale) about . Obviously, it only matters at very high energies, very short distances or very tiny time scales after the Big Bang. Then,

For a RD Universe, and thus

During the inflation, the comoving Hubble radius had to decrease by, at least, 28 orders of magnitude. The most common way to build such inflationary model is to build a model where and

It gives and , where is the time when inflation ends. In fact, we obtain that

and thus can be guessed from the so-called “e-folds”, or exponential factors, in a very simple way: note that !!!!!!!!! Therefore, more than 64 “e-folds” (or the 64th power of the number “e”) provide the necessary 10-fold we were searching for.

**Remark:** The comoving horizon is very similar to an effective time parameter!

## NEGATIVE PRESSURE

In order to allow inflation, we require that the “weak energy condition” be violated. That is, we ask that (for inflation to be possible)

It implies that or that

Obviously, there is no particle/field in the Standard Model (at least until the discovery of the Higgs field in 2012) able to do that. So, if inflation happened, it is an unknown field/particle responsible for it! The simplest implementation of inflation uses, as we said before, some scalar field . Let us define a scalar field . Its lagrangian is generally given by

The energy-momentum tensor for this scalar field can be easily obtained by the well-known prescription

Thus, we get

Negative pressure is obtained whenever we have

Therefore:

1st. A scalar field with negative pressure is trapped into a “false vacuum”.

2nd. A scalar field “slow-rolling” toward its true vacuum provides a simple model for inflation.

The evolution of the scalar field in the expanding Universe is given by a simple equation

Models where the scalar field “slow-rolls” provide (slow varying in time as a scalar function!). The time variable implies that

and it becomes negative during the inflationary phase! The slow-roll parameters

where the dot means derivative with respect to . Moreover, we also get

**Remark:** for inflation and for RD Universes imply that is a definition of inflationary phase!

## Quintessence and phantom energy

For any scalar field, and the pressure and energy density given above, we can calculate the quantity:

In the case that the kinetic term is negligible, we obtain the dark energy fluid value . However, this equation is much more general. If we have a slow-rolling scalar field over cosmological times (a very slowly time dependent scalar field) it could mimic the cosmological constant behaviour ( we are ignoring some technical problems here). If the kinetic term is NOT negligible, the value can differ from the standard value of (dark energy/cosmological constant/vacuum energy). However, current observation support a pure cosmological constant term. Moreover, this general scalar field is commonly referred as **“quintessence”** if and as **“phantom energy”** if . Both, dark energy, quintessence or phantom energy models can affect to the future of the Universe. Instead of a Big Freeze (the thermal death of the Universe), the scalar dominated Universe can even destroy atoms/matter/galaxies at some point in the future (at least on very general grounds) and terminate the Universe in a Big (or Little) Rip. Thus, let me review the possible destinies of the Universe according to modern Cosmology:

**1) Big Crunch** (or recollapse of the Universe). The Universe recollapses until the initial singularity after some time, if it is massive enough. Current observations don’t favour this case.

**2) Big Freeze** (or thermal death of the Universe). The Universe expands forever cooling itself until it reaches a temperature close to the absolute zero. It was believed that it was the only possible option with the given curvature until the discovery of dark energy in 1998.

**3) Big Rip** (or Little Rip, depending on the nature of the scalar field and the concrete model). Vacuum energy expands the Universe with an increasing rate until it “rips” even fundamental particles/atoms/matter and galaxies apart from eath other. It is a new possibility due to the existence of scalar fields and/or dark energy, a mysterious energy that makes the Universe expanss with an increasing rate overseding the gravitational pull of galaxies and clusters!

A more exotic option is that the Universe suffers “oscillations” of positively accelerated expansion and negatively accelerated expansion (oscillatory/cyclic ethernal Universes)…But that is another story…

See you in my final basic cosmological post!

# LOG#110. Basic Cosmology (V).

**Posted:**2013/06/23

**Filed under:**Cosmology, Physmatics |

**Tags:**Boltzmann equation, Compton scattering, Coulomb scattering, dark matter, free electron fraction evolution, LSP, photon decoupling, recombination, WIMP, WISP Leave a comment

## Recombination

When the Universe cooled up to , the neutrinos decoupled from the primordial plasma (soup). Protons, electrons and photons remained tighly coupled by 2 main types of scattering processes:

1) Compton scattering:

2) Coulomb scattering:

Then, there were little hydrogen (H) and though due to small baryon fraction .

The evolution of the free electron fraction provided the ratio

where and the second equality is due to the neutrality of our universe, i.e., to the fact that (by charge conservation). If remains in the thermal equilibrium, then

where we have

It gives

and the last equality is due to the fact we take . It means that at . As we have , we are out of the thermal equilibrium.

From the Boltzmann equation, we also get

or equivalently

i.e.

Using that and , we obtain

with

, the ionization rate, and

the so-called **recombination rate. **It is taken the recombination to the n=2 state of the neutral hydrogen. Note that the ground state recombination is NOT relevant here since it produces an ionizing photon, which ionizes a neutral atom, and thus the net effect is zero. In fact, the above equations provide

The numerical integration produces the following qualitative figure

The decoupling of photons from the primordial plasma is explained as

Mathematicaly speaking, this fact implies that

where is the Thomson cross section. For the processes we are interesting in, it gives

and then

Thus, we deduce that

and where implies that the decoupling of photons occurs during the time of recombination! In fact, the decoupling of photons at time of recombination is what we observe when we look at the Cosmic Microwave Background (CMB). Fascinating, isn’t it?

## Dark Matter (DM)

Today, we have strong evidences and hints that non-baryonic dark matter (DM) exists (otherwise, we should modify newtonian dynamics and or the gravitational law at large scales, but it seems that even if we do that, we require this dark matter stuff).

In fact, from cosmological observations (and some astrotronomical and astrophysical measurements) we get the value of the DM energy density

The most plausible candidate for DM are the Weakly Interacting Massive Particles (WIMPs, for short). Generic WIMP scenarios provide annihilations

where is some “heavy” DM particle and the (ultra)weak interaction above produces light particles in form of leptons and antileptons, tighly couple to the cosmic plasma. The Boltzmann equation gives

Define the yield (or ratio) . It is produced since generally we have

and since , then . Thus,

and

Now, we can introduce a new time variable, say

Then, we calculate

For a radiation dominated (RD) Universe, implies that and

In this case, we obtain

with

The final freeze out abundance is got in the limit . Typically, , and when and , for , and there, the yield drops exponentially

or

Integrating this equation,

and then

Generally, and the freeze out temperature for WIMPs is got with the aid of the following equation

Indeed,

A qualitative numerical solution of the “WIMP” miracle (and its freeze out) is given by the following sketch

The present abundance of heavy particle relics gives

and where the effect of entrpy dumping after the freeze-out is encoded into the factor

with

Moreover, the DM energy density can also be estimated:

so

The main (current favourite) candidate for WIMP particles are the so called lightest supersymmetric particles (LSP). However, there are other possible elections. For instance, Majorana neutrinos (or other sterile neutrino species), Z prime bosons, and other exotic particles. We observe that here there is a deep connection between particle physics, astrophysics and cosmology when we talk about the energy density and its total composition, from a fundamental viewpoint.

**Remark:** there are also WISP particles (Weakly Interacting Slim Particles), like (superlight) axions and other exotics that could contribute to the DM energy density and/or the “dark energy”/vacum energy that we observe today. There are many experiments searching for these particles in laboratories, colliders, DM detection experiments and astrophysical/cosmological observations (cosmic rays and other HEP phenomena are also investigated to that goal).

**See you in a next cosmological post!**

# LOG#108. Basic Cosmology (III).

**Posted:**2013/06/09

**Filed under:**Cosmology, Physmatics |

**Tags:**annihilation rate, Boltzmann equation, Bose-Einstein distribution, Cosmology, cross-section, decoupling, equilibrium, equilibrium rate, Fermi-Dirac distribution, freeze out, Maxwell-Boltzmann distribution, out-of-equilibrium processes in cosmology, Saha equation, thermal equilibrium 1 Comment

The current Universe has evolved since its early phase of thermal equilibrium until the present state. The departure from thermal equilibrium in the early Universe made a fossil record we can observe at current time!

There are some easy rules for thermal equilibrium. The easiest one, is that coming from the “interaction rate” . It can be expressed in the following way:

and then, at a given temperature T, we get

and where

**Remark:** If , then

and it implies that a particle interacts less than once after the time .

Moreover, we can understand roughly the so-called decoupling era:

**1st. Any interaction mediated by a massless gauge boson provides**

with and

and this implies that

and

so the equilibrium temperature is found whenever !

**2nd. Interactions mediated by any massive gauge boson provides**

with

and this implies that

and

and then

Moreover,

As a consequence, we can realize that the out-of-equilibrium phenomena in the early and current Universe are very important processes! In particular:

1) They provide the formation of light elements during the * Big Bang Nucleosynthesis (BBN)*, also known as

**, i.e., the formation of the first light elements after the Big Bang (circa 300000 years after the Universe “birth”).**

*primordial nucleosynthesis*2) They provide the path of * recombination* of electrons and protons into hydrogen atoms.

3) They imply the (likely) production of dark matter (or equivalently the presence of some kind of “modified gravity” or/and modified newtonian dynamics).

**Boltzmann’s equation for annihilation of particles in equilibrium**

There is a beautiful equation that condenses the previous physical process of equilibrium at a given temperature and the particle production it yields. Conceptually speaking, we have

Consider a process like

and where the particle 1 is the one we are interested in. Then, we deduce that

where A is certain complicated facter involving “delta functions” of the energies and momenta of the particles 1,2,3,4 and an additional term depending on the statistics of the particle. Explicitly, it takes the form

with

and where

is the Fermi-Dirac (FD, -)/Bose-Einstein (BE,+) distribution. In fact, the above FD/BE factors provide the so-called Pauli blocking/”Bose-Einstein” enhancement effects for the particle production in the processes and . Indeed, particle physics enter into the game here (see above formulae again) and we assume

Do you recognize the *principle* of *detailed balance *in this equation?

We can simplify the assumptions a little bit:

1st. The kinetic equilibrium is taken to be a rapid elastic scattering and we input the FD/BE statistics without loss of generality.

2nd. The annihilation in thermal equilibrium will be calculated from the sum of the chemical potential in any balanced equation.

3rd. Low temperature approximation. Suppose that

then we obtain the Maxwell-Boltzmann approximation to the FD/BE statistics

and , so, since , we get

What is ? After a change of variable

is the “equilibrium number density” deducen from the expression

and it yields

and then we finally get that

equals

Now, we can define the thermally averaged cross section

The Boltzmann equation becomes with these conventions

Remark (I): LHS is similar to and the RHS is similar to

Remark (II): If the reaction rate is , then it provides the chemical equilibrium condition well known in the nuclear statistical equilibrium as the Saha equation, i.e.,

# LOG#106. Basic Cosmology (I).

**Posted:**2013/05/26

**Filed under:**Cosmology, General Relativity, Physmatics |

**Tags:**Big Bang, Bose-Einstein distribution, cosmic microwave background, Cosmological principle, Cosmology, curvature parameter, curved Universe, dark energy, degrees of freedom, dust, early Universe, Einstein tensor, Einstein-Hilbert action, energy density, energy-momentum tensor, equivalence principle, Fermi-Dirac distribution, General Relativity, geodesic equation, geodesics, hot ideal gas, ideal gas, Killing equation, Killing vector, maximally symmetric space, natural units, neutrinos, number density, parsec, particle physics, perfect cosmological principle, perfect fluid, plane Universe, pressure, redshift, relativistic matter, Standard Cosmological Model, thermal equilibrium, yield Leave a comment

The next thread is devoted to Cosmology. I will intend to be clear and simple about equations and principles of current Cosmology with a General Relativity background.

First of all…I will review the basic concepts of natural units I am going to use here. We will be using the following natural units:

We will take the Planck mass to be given by

The solar mass is and the parsec is given by the value

Well, current Cosmology is based on General Relativity. Even if I have not reviewed this theory with detail in this blog, the nice thing is that most of Cosmology can be learned with only a very little knowledge of this fenomenal theory. The most important ideas are: metric field, geodesics, Einstein equations and no much more…

In fact, newtonian gravity is a good approximation in some particular cases! And we do know that even in this pre-relativistic theory

via the Poisson’s equation

This idea, due to the equivalence principle, is generalized a little bit in the general relativistic framework

The spacetime geometry is determined by the metric tensor . The matter content is given by the stress-energy-momentum tensor . As we know one of these two elements, we can know, via Eisntein’s field equations the another. That is, given a metric tensor, we can tell how energy-momentum “moves” in space-time. Given the energy-momentum tensor, we can know what is the metric tensor in spacetime and we can guess how the spacetime bends… This is the origin of the famous motto: “Spacetime says matter how to move, energy-momentum says spacetime how to curve”! Remember that we have “deduced” the Einstein’s field equations in the previous post. Without a cosmological constant term, we get

Given a spacetime metric , we can calculate the (affine/Levi-Civita) connection

The Riemann tensor that measures the spacetime curvature is provided by the equation

The Ricci tensor is defined to be the following “trace” of the Riemann tensor

The Einstein tensor is related to the above tensors in the well-known manner

The Einstein’s equations can be derived from the Einstein-Hilbert action we learned in the previous post, using the action principle and the integral

The geodesic equation is the path of a freely falling particle. It gives a “condensation” of the Einstein’s equivalence principle too and it is also a generalization of Newton’s law of “no force”. That is, the geodesic equation is the feynmanity

Finally, an important concept in General Relativity is that of isometry. The symmetry of the “spacetime manifold” is provided by a Killing vector that preserves transformations (isometries) of that manifold. Mathematically speaking, the Killing vector fields satisfy certain equation called the Killing equation

Maximally symmetric spaces have Killing vectors in n-dimensional (nD) spacetime. There are 3 main classes or types of 2D maximally symmetric that can be generalized to higher dimensions:

1. The euclidean plane .

2. The pseudo-sphere . This is a certain “hyperbolic” space.

3. The spehre . This is a certain “elliptic” space.

**The Friedmann-Robertson-Walker Cosmology**

Current cosmological models are based in General Relativity AND a simplification of the possible metrics due to the so-called Copernican (or cosmological) principle: the Universe is pretty much the same “everywhere” you are in the whole Universe! Remarkbly, the old “perfect” Copernican (cosmological) principle that states that the Universe is the same “everywhere” and “every time” is wrong. Phenomenologically, we have found that the Universe has evolved and it evolves, so the Universe was “different” when it was “young”. Therefore, the perfect cosmological principle is flawed. In fact, this experimental fact allows us to neglect some old theories like the “stationary state” and many other “crazy theories”.

What are the observational facts to keep the Copernican principle? It seems that:

1st. The distribution of matter (mainly galaxies, clusters,…) and radiation (the cosmic microwave background/CMB) in the observable Universe is **homogenous and isotropic.**

2nd. The Universe is NOT static. From Hubble’s pioneer works/observations, we do know that galaxies are receeding from us!

Therefore, these observations imply that our “local” Hubble volume during the Hubble time is similar to some spacetime with homogenous and isotropic spatial sections, i.e., it is a spacetime manifold . Here, denotes the time “slice” and represents a 3D maximally symmetric space.

The geometry of a locally isotropic and homogeneous Universe is represented by the so-called Friedmann-Robertson-Walker metric

Here, is the called the **scale factor. **The parameter determines the geometry type (plane, hyperbolic or elliptical/spherical):

1) If , then the Universe is “flat”. The manifold is .

2) If , then the Universe is “open”/hyperbolic. The manifold would be .

3) If , then the Universe is “closed”/spherical or elliptical. The manifold is then .

**Remark:** The ansatz of local homogeneity and istoropy only implies that the spatial metric is locally one of the above three spaces, i.e., . It could be possible that these 3 spaces had different global (likely topological) properties beyond these two properties.

**Kinematical features of a FRW Universe**

The first property we are interested in Cosmology/Astrophysics is “distance”. Measuring distance in a expanding Universe like a FRW metric is “tricky”. There are several notions of “useful” distances. They can be measured by different methods/approaches and they provide something called sometimes “the cosmologidal distance ladder”:

1st. **Comoving distance.** It is a measure in which the distance is “taken” by a fixed coordinate.

2nd. **Physical distance.** It is essentially the comoving distance times the scale factor.

3rd.** Luminosity distance.** It uses the light emitted by some object to calculate its distance (provided the speed of light is taken constant, i.e., special relativity holds and we have a constant speed of light)

4th. **Angular diameter distance. **Another measure of distance using the notion of parallax and the “extension” of the physical object we measure somehow.

There is an important (complementary) idea in FRW Cosmology: the** particle horizon**. Consider a light-like particle with . Then,

or

The total comoving distance that light have traveled since a time is equal to

It shows that NO information could have propagated further and thus, there is a “comoving horizon” with every light-like particle! Here, this time is generally used as a “conformal time” as a convenient tiem variable for the particle. The physical distance to the particle horizon can be calculated

There are some important kinematical equations to be known

A) **For the geodesic equation, the free falling particle,** we have

for the FRW metric and, moreover, the energy-momentum vector is defined by the usual invariant equation

This definition defines, in fact, the proper “time” implicitely, since

and the 0th component of the geodesic equation becomes

Therefore we have deduced that . This is, in fact, the socalled “redshift”. The cosmological redshift parameter is more generally defined through the equation

B) **The Hubble’s law.**

The luminosity distance measures the flux of light from a distant object of known luminosity (if it is not expanding). The flux and luminosity distance are bound into a single equation

If we use the comoving distance between a distant emitter and us, we get

for a expanding Universe! That is, we have used the fact that luminosity itself goes through a comoving spherical shell of radius . Moreover, it shows that

The luminosity distance in the expanding shell is

and this is what we MEASURE in Astrophysics/Cosmology. Knowing , we can express the luminosity distance in terms of the redshift. Taylor expansion provides something like this:

where higher order terms are sometimes referred as “statefinder parameters/variables”. In particular, we have

and

C) **Angular diameter distance.**

If we know that some object has a known length , and it gives some angular “aperture” or separation , the angular diameter distance is given by

The comoving size is defined as , and the coming distance is again . For “flat” space, we obtain that

that is

In the case of “curved” spaces, we get

**FRW dynamics**

Gravity in General Relativity, a misnomer for the (locally) relativistic theory of gravitation, is described by a metric field, i.e., by a second range tensor (covariant tensor if we are purist with the nature of components). The metric field is related to the matter-energy-momentum content through the Einstein’s equations

The left-handed side can be calculated for a FRW Universe as follows

The right-handed side is the energy-momentum of the Universe. In order to be fully consistent with the symmetries of the metric, the energy-momentum tensor MUST be diagonal and . In fact, this type of tensor describes a perfect fluid with

Here, are functions of (cosmological time) only. They are “state variables” somehow. Moreover, we have

for the fluid at rest in the comoving frame. The Friedmann equations are indeed the EFE for a FRW metric Universe

for the 00th compoent as “constraint equation.

for the iith components.

Moreover, we also have

and this conservation law implies that

Therefore, we have got two independent equations for three unknowns . We need an additional equation. In fact, the equation of state for provides such an additional equation. It gives the “dynamics of matter”!

In summary, the basic equations for Cosmology in a FRW metric, via EFE, are the Friedmann’s equations (they are secretly the EFE for the FRW metric) supplemented with the energy-momentum conservations law and the equation of state for the pressure :

1)

2)

3)

There are many kinds of “matter-energy” content of our interest in Cosmology. Some of them can be described by a simple equation of state:

Energy-momentum conservation implies that . 3 special cases are used often:

1st. **Radiation (relativistic “matter”).** and thus, and

2nd. **Dust (non-relativistic matter).** . Then, and

3rd.** Vacuum energy (cosmological constant).** . Then, and

**Remark (I):** Particle physics enters Cosmology here! Matter dynamics or matter fields ARE the matter content of the Universe.

**Remark (II):** Existence of a Big Bang (and a spacetime singularity). Using the Friedmann’s equation

if we have that , the so-called weak energy condition, then should have been reached at some finite time in the past! That is the “Big Bang” and EFE are “singular” there. There is no scape in the framework of GR. Thus, we need a quantum theory of gravity to solve this problem OR give up the FRW metric at the very early Universe by some other type of metric or structure.

**Particles and the chemical equilibrium of the early Universe**

Today, we have DIRECT evidence for the existence of a “thermal” equilibrium in the early Universe: the cosmic microwave background (CMB). The CMB is an isotropic, accurate and non-homogeneous (over certain scales) blackbody spectrum about !

Then, we know that the early Universe was filled with a hot dieal gas in thermal equilibrium (a temperature can be defined there) such as the energy density and pressure can be written in terms of this temperature. This temperature generates a distribution . The number of phase space elements in is

and where the RHS is due to the uncertainty principle. Using homogeneity, we get that, indeed, , and where we can write the volume . The energy density and the pressure are given by (natural units are used)

When we are in the thermal equilibrium at temperature T, we have the Bose-Einstein/Fermi-Dirac distribution

and where the is for the Fermi-Dirac distribution (particles) and the is for the Bose-Einstein distribution (particles). The number density, the energy density and the pressure are the following integrals

And now, we find some special cases of matter-energy for the above variables:

1st. **Relativistic, non-degenerate matte**r (e.g. the known neutrino species). It means that and . Thus,

2nd. **Non-relativistic matter** with only. Then,

, and

The total energy density is a very important quantity.** In the thermal equilibrium,** the energy density of non-relativistic species is exponentially smaller (suppressed) than that of the relativistic particles! In fact,

for radiation with

and the effective degrees of freedom are

**Remark:** The factor in the DOF and the variables above is due to the relation between the Bose-Einstein and the Fermi-Dirac integral in d=3 space dimensions. In general d, the factor would be

**Entropy conservation and the early Universe**

The entropy in a comoving volume IS a conserved quantity IN THE THERMAL EQUILIBRIUM. Therefore, we have that

and then

or

Now, since

then

if we multiply by and use the chain rule for , we obtain

but it means that , where is the entropy density defined by

Well, the fact is that we know that the entropy or more precisely the entropy density is the early Universe is dominated by relativistic particles ( this is “common knowledge” in the Stantard Cosmological Model, also called ). Thus,

It implies the evolution of temperature with the redshift in the following way:

Indeed, since we have that , , the **yield** variable

is a convenient quantity that represents the “abundance” of decoupled particles.

See you in my next cosmological post!