which is predominantly made up of the active left-handed SM neutrinoνL. This is called thetype I seesaw mechanism (cf. above) and it is now possible to get the right masses for the SM neutrinos with moderate values of the neutrino Yukawa couplings (of the same order of magnitude as, say, the electron Yukawa coupling).
Here, we have indicated to which factor of GPS the corresponding representations belong. The Pati-Salam gauge bosons in turn contain the MSSM gauge bosons as described in section 2.2.2. The gauge bosons residing in the (6,2,2), on the other hand, get GUT-scale masses in the breaking ofSO(10) toGPS.
• MSSM Higgses
The (1,2,2)-representation of GPS, which contains both the MSSM up-type Higgs and down-type Higgs doublets, fits into the following representation ofSO(10) 9
h10∼10= (1,2,2)
| {z } h
⊕(6,1,1). (2.38)
Here a potential problem of SO(10) unification emerges: the additional degrees of freedom contained in the (6,1,1) include colour triplets under GSM, which would lead to rapid proton decay. This can only be prevented by making these colour-triplets very heavy (which suppresses the dangerous operators responsible for proton decay), while keeping the MSSM Higgses sufficiently light. This problem is referred to as thedoublet-triplet splitting problem in the literature and as already mentioned above we assume that a solution to this problem is provided in the first stage of the breaking chain Eqn. (2.35). A more detailed discussion of the doublet-triplet splitting problem within supersymmetric GUTs and possible solutions (like e.g. the Dimopoulos-Wilczek mechanism [91], cf. also [92]) can for example be found in [74, 93].
With these definitions and sticking to renormalisable operators, the allowed terms in the superpotential read
W =yF.h10.F +1
2µh10.h10. (2.39)
After the breaking to GPS and then to GSM this superpotential reproduces the Yukawa couplings and the µ-term, cf. Eqns. (2.26) and (2.3).
We can also embed the Higgs fields responsible for the breaking ofGPStoGSMinto two additional spinor representations16and16. We call these representations¯ H16 andH¯16, respectively. Given that they have to get VEVs in the right-handed neutrino direction in order to break to the right SM vacuum, operators of the form
λ
ΛF.H¯16F.H¯16, (2.40)
again yield a Majorana mass term for the right-handed neutrinos, which together with the superpotential Eqn. (2.39) gives small masses to the active SM neutrinos through the type I Seesaw mechanism. Since these same operators also play a crucial role in the models of inflation we are going to study in chapter 7, they can provide a link between particle
9 Remember that for SU(2),2and ¯2are equivalent. Here, we use the fundamental representation2 for the SU(2)L factor for convenience, in contrast to section 2.2.1, where we used the anti-fundamental representation¯2.
Name SO(10) GPS Decomp. Contains MSSM Fields Sector
F 16 (4,2,1) q,l Matter
-(¯4,1,¯2) uc,dc,lc,νc (Super)Fields A 45 (15,1,1) Ga,A15
(1,1,3) A18 Gauge Bosons &
(1,3,1) W+,W0,W− Gauginos (6,2,2)
h10 10 (1,2,2) hu,hd Higgs Bosons &
(6,1,1) Higgsinos
Table 2.3: Decomposition of some SO(10) representations with respect to GPS and the MSSM fields that are embedded within these representations. For simplicity we only show the first family of matter fields. Colour indices are suppressed.
physics on the one hand and early universe cosmology on the other hand. The study of such possible connections is one of the main subjects of this thesis.
In section 7.1.1, we explicitly show how such operators can be generated by integrating out heavy messenger particles.
Cosmological Inflation
After having introduced the necessary particle physics tools and models in the previous two chapters we now turn our attention to cosmology. In particular, we are going to discuss two very important events in the history of our universe: cosmological inflation and the generation of the baryon asymmetry. The first, inflation, was originally introduced by Starobinsky and Guth [7] and later in a modified version (called slow-roll inflation) also by Linde, Albrecht and Steinhardt [12, 13] to solve what is commonly referred to as the horizon and flatness problems of the Standard Hot Big Bang Model (SBB). It has since developed into an integral part of our understanding of the early universe, as inflation can, for example, also account for the origin of the large-scale structure we observe in the universe today. The second, baryogenesis, addresses the question of how we can live in a universe filled with matter but no appreciable amount of antimatter given that our universe started out in a state that is symmetric between the two.
In this chapter we focus on inflation, starting with a discussion of the issues that motivated its introduction and then discussing how inflation works in its most popular manifestation, slow-roll inflation. After that we survey the zoology of inflationary models and address some of the problems connected to the embedding of inflationary models within supergravity.
The motivation for inflation, the basics of slow-roll inflation and the different types of inflationary models are discussed in a large number of review articles and text books. The ones we have mainly consulted include [94, 95, 96, 97, 98, 99, 100, 101, 102, 103, 104].
Baryogenesis and in particular baryogenesis through leptogenesis are discussed in the next chapter.
3.1 The Problem of Initial Conditions
First, we want to discuss in more detail the problems that led to the idea of inflation and sketch the way inflation actually solves these problems. When talking about the Standard Hot Big Bang Model in this section, we always mean a matter or radiation dominated universe, that expands and cools as it evolves, but doesnot include a stage of accelerated expansion as part of its history.
The Horizon Problem
The first of these problems is called thehorizon problem and it is connected to the homo-geneity and isotropy of the Cosmic Microwave Background Radiation (CMBR) that we can observe today and that fills up the entire sky.
This radiation, which has a temperature of [105]
TCMB= 2.72548±0.00057 K (3.1)
and exhibits a nearly perfect black body spectrum, is incredibly homogeneous and isotropic [106, 6]. This poses a serious problem for the SBB model, since it turns out that in a universe with a normal thermal history without a phase of accelerated expansion, the radiation reaching us today from different directions in the sky must have originated from regions of spacetime that could never have been in causal contact at the time the radiation was emitted. How, then, can it be so extremely homogeneous and isotropic? Let us quantify this problem to show its severity.
The CMB radiation was created everywhere in the universe in an event commonly termed last scattering at a redshift of
1 +zls≡ a(t0)
a(tls) ≈1090, (3.2)
approximately 378000 years after the Big Bang [6, 66]. However, the radiation we receive at earth today originated from a spherical surface around the earth, called the surface of last scattering. The radius rls of this surface increased from the time of last scattering until today, following the expansion of the universe
rls∝a(t). (3.3)
Since the CMBR is visible to us today, the current radius of this surface must be smaller that the radius of the currently visible universe, roughly given by the Hubble distance
rvisible' H−1. (3.4)
However, for a universe filled with “ordinary matter” described by an equation of state with 1 + 3ω >0 (cf. Appendix C for more details), the horizon distance rvisible expands faster than the distance between two comoving points following the expansion of the universe.
Thus, more and more parts of the universe “fall inside our horizon” as the universe evolves.
In particular
rvisible∝t , rls∝tf,
(3.5) (3.6) with 0 < f < 1 for matter or radiation dominance. Going back in time, on the other hand, this means that what is visible to us today, and thus is in causal contact, must not necessarily have been in causal contact in the past. This is in particular true for different
regions on the surface of last scattering. Extrapolating back what we know about our universe – its age, its energy content and the time of last scattering – we find that at the creation of the CMB radiation, the surface of last scattering consisted of a huge number of causally disconnected regions. How many? Viewed from earth, the horizon size at the time of last scattering subtended only an angle of about 1 degree on the surface of last scattering [102].
The solution inflation provides to this problem is that inflation can stretch a causal patch that initially lay well inside the horizon to a size much large than the horizon, large enough to encompass the whole surface of last scattering and thereby establishing causal contact on it 1.
The Flatness Problem
The second problem is connected to the geometry of our universe. We know from obser-vation that today our universe is almost perfectly spatially flat [6]
Ω0= 1.02±0.02, (3.7)
cf. Eqn. (C.19). However, from the Friedmann equation (C.18) we get d|Ω−1|
d lna >0 for 1 + 3ω >0, (3.8) which holds, in particular, for a matter or radiation dominated universe. Extrapolating back e.g. to the time of last scattering or Big Bang Nucleosynthesis (BBN), respectively, the geometry of our universe must have been incredibly fine-tuned [102],
|Ωls−1| <0.0004,
|ΩBBN−1| <10−12,
(3.9) (3.10) to end up with the universe we see today. Why should the universe have started out in such a very special state?
An answer to thisflatness problemis again given by a phase of accelerated expansion of the universe, because during such a phase the geometry of our universe is actually driven towards flatness, not away from it. This can again be seen directly from the Friedmann equation (C.18), because during inflation the comoving Hubble radius (aH)−1 decreases2, such that any “curvature term” on the right hand side of Eqn. (C.18) becomes negligible.
Topological Defects
Another problem that inflation can solve is the problem of topological defects (also called the monopole problem on occasion, because monopoles turn out to be particularly danger-ous). These topological defects can be generated e.g. during phase transitions in the early
1 As we will see shortly, during inflation the physical horizon size stays almost constant whereas the physical distance between comoving observers grows exponentially, cf. Eqns. (3.20).
2 During inflation ¨a >0 by definition. Thus, as ˙a >0 increases, (aH)−1= ˙a−1 decreases.
universe when a GUT symmetry is broken down spontaneously to a smaller symmetry group [18, 107].
Whether or not and what kinds of topological defects are generated in such a symme-try breaking phase transition depends on group-theoretical issues 3 as well as the exact mechanism that governs the phase transition 4.
In most of the standard GUT scenarios, however, these topological defects are quite copiously produced and an estimate shows that they would give rise to a catastrophically large contribution to the energy density of the universe Ωtop 1 (cf. for example [96]
for a detailed discussion). This would dramatically over-close the universe and make it re-collapse onto itself, which seems to rule out most of the GUT scenarios.
If, however, these phase transitions take place before inflation, then during inflation the energy density contribution of the produced defects can be diluted to an acceptably low value Ωtop 1. For GUT model building, on the other hand, this means that any phase transition that happens after inflation must not produce additional topological defects5. For a more detailed discussion of topological defects in the early universe we refer the reader to [111, 96, 100] and references therein.
Initial Perturbations
The last point we want to discuss as a motivation for inflation is the question what set the initial conditions for the structure formation in our universe. Whereas it might seem that after inflation the universe is a pretty boring place, void of any structure, this is actually not the case. While it is true that inflation tends to wash out any localised structure or energy density perturbation present before the onset of inflation, at the same time it provides its own mechanism to generate such energy density perturbations: during inflation the quantum fluctuations of the inflaton field can get stretched across the horizon and become classical quantities [8]. These field perturbations are coupled to the spacetime metric, leading to gravitational perturbations of the homogeneous background metric, which in turn lead to temperature anisotropies in the CMB radiation and later on to the formation of the large scale structure in our universe.
The fact that these perturbations can be computed for a given inflationary model and be compared to experimental data on the temperature anisotropies of the CMB, showing an astonishing agreement between theory and experiment for a wide class of inflationary models, firmly established inflation as part of our understanding of early universe cosmology. As a matter of fact, inflation is the only known mechanism that can naturally reproduce the peak structure of the CMB power spectrum (and simultaneously
3 To be more precise, it depends on the fact whether or not one of the homotopy groupsπi(M) of the vacuum manifoldMis non-trivial. The vacuum manifold is given by the quotient groupM=G/H, where G the group that is spontaneously broken down toH⊂G.
4In section 7.1 we construct a model that avoids the production of topological defects in such a phase transition, notwithstanding the fact that they would be allowed from a group-theoretical perspective.
5 Cosmic strings might be allowed to a certain extend and it was even thought that they might help to push the preferred value of the spectral index ns towards (or even above)ns = 1, cf. e.g. [108, 109].
According to a more recent analysis [110], however, a scale invariant initial perturbation spectrum with ns= 1 is now disfavoured at 2.4σeven if strings are present.
solve the horizon, flatness, and monopole problems). This comes about because once the fluctuation modes have been stretched across the horizon and become classical quantities, the all re-enter the horizon in phase and their interference pattern produces the peak structure in the CMB power spectrum. A more thorough discussion of this very interesting topic can be found in [112, 104].
In the remainder of this chapter, we discuss these points in more detail.