5.5. Slowly-Rolling Scalar Fields
If inflation is to solve the problems of the standard cosmology, then it must be active at extremely early times. Thus, we would like to address the earliest times in the universe amenable to a classical description. We expect this to be at or around the Planck time tP and since Planckian quantities arise often in inflation we will retain values of the Planck mass in the equations of this section. There are many models of inflation, but because of time constraints we will concentrate almost exclusively on the chaotic inflation model of Linde. We have borrowed heavily in places here from the excellent text of Liddle and Lyth [7].
Consider modeling matter in the early universe by a real scalar field
, with potential
V(
).
The energy-momentum tensor for
is
![]() |
(188) |
For simplicity we will specialize to the homogeneous case, in which all quantities depend only on cosmological time t and set k = 0. A homogeneous real scalar field behaves as a perfect fluid with
![]() |
(189)
|
The equation of motion for the scalar field is given by
![]() |
(191) |
which can be thought of as the usual equation of motion for a scalar field in Minkowski space, but with a friction term due to the expansion of the universe. The Friedmann equation with such a field as the sole energy source is
![]() |
(192) |
A very specific way in which accelerated expansion can occur is if the
universe is dominated by an energy component that approximates a
cosmological constant. In that case the associated expansion rate will
be exponential, as we have already seen. Scalar fields can accomplish
this in an interesting way. From (189) it is clear that if
2 <<
V(
) then
the potential energy of the
scalar field is the dominant contribution to both the energy density
and the pressure, and the resulting equation of state is
p
-
, approximately
that of a cosmological constant. the resulting
expansion is certainly accelerating. In a loose sense, this negligible
kinetic energy is equivalent to the fields slowly rolling down its
potential; an approximation which we will now make more formal.
Technically, the slow-roll approximation for inflation involves
neglecting the
term
in (191) and neglecting the kinetic energy of
compared to
the potential energy. The
scalar field equation of motion and the Friedmann equation then become
![]() |
(193) |
![]() |
(194) |
where in this lecture a prime denotes a derivative with respect to
.
These conditions will hold if the two slow-roll conditions are satisfied. These are
![]() |
(195) |
where the slow-roll parameters are given by
![]() |
(196) |
and
![]() |
(197) |
It is easy to see that the slow roll conditions yield
inflation. Recall that inflation is defined by
/ a > 0. We can
write
![]() |
(198) |
so that inflation occurs if
![]() |
(199) |
But in slow-roll
![]() |
(200) |
which will be small. Smallness of the other parameter
helps
to ensure that inflation will continue for a sufficient period.
It is useful to have a general expression to describe how much inflation occurs, once it has begun. This is typically quantified by the number of e-folds, defined by
![]() |
(201) |
Usually we are interested in how many efolds occur between a given
field value
and the field value at the end of inflation
end,
defined by
(
end) =
1. We also
would like to express N in terms of the potential. Fortunately this
is simple to do via
![]() |
(202) |
The issue of initial conditions for inflation is one that is quite subtle and we will not get into a discussion of that here. Instead we will remain focused on chaotic inflation, in which we assume that the early universe emerges from the Planck epoch with the scalar field taking different values in different parts of the universe, with typically Planckian energies. There will then be some probability for inflation to begin in some places, and we shall focus on those.