It is rather frustrating that the only component of the universe which
we understand theoretically is the radiation! While understanding the
baryonic and dark matter components [in particular the values of
B and
DM] is by
no means trivial, the issue of dark
energy is lot more perplexing, thereby justifying the attention it has
received recently. In this section we will discuss several aspects of
the dark energy problem.
The key observational feature of dark energy is that - treated as a
fluid with a stress tensor Tba =
dia (,
-p, -p, -p)
- it has an equation state p =
w
with
w
-0.8 at the
present epoch. The spatial part g of the geodesic acceleration
(which measures the
relative acceleration of two geodesics in the spacetime) satisfies an
exact equation in general relativity given by:
![]() |
(190) |
This shows that the source of geodesic acceleration is
( + 3p) and
not
.
As long as (
+
3p) > 0, gravity remains attractive while
(
+
3p) < 0 can
lead to repulsive gravitational effects. In other words, dark energy
with sufficiently negative pressure will
accelerate the expansion of the universe, once it starts dominating
over the normal matter. This is precisely what is established from the
study of high redshift supernova, which can be used to determine the
expansion rate of the universe in the past
[6].
The simplest model for a fluid with negative pressure is the
cosmological constant (for some recent reviews, see
[23])
with w = -1,
= -p = constant.
If the dark energy is indeed a cosmological constant, then it introduces
a fundamental length scale in the theory
L
H
-1, related to the constant dark energy
density
DE by
H
2
(8
G
DE / 3).
In classical general relativity, based on the constants G,
c and
L
, it
is not possible to construct any dimensionless combination from these
constants. But when one introduces the Planck constant,
, it is
possible to form the dimensionless combination
H
2(G
/
c3)
(Lp2 /
L
2).
Observations then require (Lp2 /
L
2)
10-123.
As has been mentioned several times in literature, this will require
enormous fine tuning. What is more, in the past, the energy density of
normal matter and radiation would have been higher while the energy
density contributed by the cosmological constant
does not change. Hence we need to adjust the energy densities
of normal matter and cosmological constant in the early epoch very
carefully so that
NR
around the current
epoch. This raises the second of the two cosmological constant problems:
Why is it that (
/
NR)
=
(1) at the
current phase of the universe ?
Because of these conceptual problems associated with the cosmological
constant, people have explored a large variety of alternative
possibilities. The most popular among them uses a scalar field
with a suitably chosen potential
V(
) so
as to make the vacuum energy vary with time. The hope then is that, one
can find a model in which the current value can be explained naturally
without any fine tuning.
A simple form of the source with variable w are scalar fields with
Lagrangians of different forms, of which we will discuss two possibilities:
![]() |
(191) |
Both these Lagrangians involve one arbitrary function
V(). The
first one, Lquin, which is a natural generalization of
the Lagrangian for a non-relativistic particle, L = (1/2)
2 -
V(q), is usually called
quintessence (for a small sample of models, see
[24];
there is an extensive
and growing literature on scalar field models and more references can be
found in the reviews in ref.
[23]).
When it acts as a source in Friedmann universe,
it is characterized by a time dependent w(t) with
![]() |
(192) |
The structure of the second Lagrangian (which arise in string theory
[25])
in Eq. (191) can be understood by a simple analogy from
special relativity. A relativistic particle with (one dimensional) position
q(t) and mass m is described by the Lagrangian
L = -m (1 -
2)1/2.
It has the energy E = m / (1 -
2)1/2
and momentum k = m
/ (1 -
2)1/2 which are related by
E2 = k2 + m2. As
is well known, this allows the possibility of having massless
particles with finite
energy for which E2 = k2. This is
achieved by taking the limit of m
0
and
1,
while keeping the ratio in E = m
/ (1 -
2)1/2
finite. The momentum acquires a life of its own, unconnected with the
velocity
,
and the energy is expressed in terms of the momentum
(rather than in terms of
) in the
Hamiltonian formulation. We can now construct a field theory by upgrading
q(t) to a field
. Relativistic
invariance now requires
to depend on
both space and time
[
=
(t,
x)] and
2 to be replaced by
i
i
. It is also
possible now to treat the mass parameter m as a function of
, say,
V(
)
thereby obtaining a field theoretic Lagrangian L =
-V(
)
(1 -
i
i
)1/2.
The Hamiltonian structure of this
theory is algebraically very similar to the special relativistic example we
started with. In particular, the theory allows solutions in which
V
0,
i
i
1 simultaneously,
keeping the energy (density) finite. Such
solutions will have finite momentum density (analogous to a massless
particle with finite momentum k) and energy density. Since the
solutions can now depend on both space and time (unlike the special
relativistic example in which
q depended only on time), the momentum density can be an arbitrary
function of the spatial coordinate. The structure of this Lagrangian is
similar to those analyzed in a wide class of models called K-essence
[26]
and provides a rich gamut of
possibilities in the context of cosmology
[27,
28].
Since the quintessence field (or the tachyonic field) has
an undetermined free function
V(), it
is possible to choose this function in order to produce a given
H(a). To see this explicitly, let
us assume that the universe has two forms of energy density with
(a) =
known(a) +
(a)
where
known(a) arises from any known forms of
source (matter, radiation,
...) and
(a) is due to a scalar field. Let us first
consider quintessence. Here, the potential is given implicitly by the form
[29,
27].
![]() |
(193) |
where Q(a)
[8
G
known(a) /
3H2(a)] and prime
denotes differentiation with respect to a. (The result used in
Eq. (155) is just a special case of this when Q = 0) Given any
H(a), Q(a), these equations determine
V(a) and
(a)
and thus the potential
V(
).
Every quintessence model studied in
the literature can be obtained from these equations.
Similar results exists for the tachyonic scalar field as well
[27].
For example, given any H(a), one can construct a
tachyonic potential
V() so
that the scalar field is the source for
the cosmology. The equations determining
V(
) are
now given by:
![]() |
(194) |
Equations (194) completely solve the problem. Given any
H(a), these equations determine V(a) and
(a) and
thus the potential
V(
).
A wide variety of phenomenological models with time dependent
cosmological constant have been considered in the literature all of
which can be mapped to a scalar field model with a suitable
V(
).
While the scalar field models enjoy considerable popularity (one reason being they are easy to construct!) it is very doubtful whether they have helped us to understand the nature of the dark energy at any deeper level. These models, viewed objectively, suffer from several shortcomings:
They completely lack predictive power. As
explicitly demonstrated above, virtually every form of
a(t) can be modeled by a suitable "designer"
V().
All the scalar field potentials require fine tuning of the parameters in order to be viable. This is obvious in the quintessence models in which adding a constant to the potential is the same as invoking a cosmological constant. So to make the quintessence models work, we first need to assume the cosmological constant is zero. These models, therefore, merely push the cosmological constant problem to another level, making it somebody else's problem!.
By and large, the potentials used in the literature have no natural field theoretical justification. All of them are non-renormalisable in the conventional sense and have to be interpreted as a low energy effective potential in an ad hoc manner.
One key difference between cosmological constant and scalar field models is that the latter lead to a w(a) which varies with time. If observations have demanded this, or even if observations have ruled out w = -1 at the present epoch, then one would have been forced to take alternative models seriously. However, all available observations are consistent with cosmological constant (w = -1) and - in fact - the possible variation of w is strongly constrained [30] as shown in Figure 6.
![]() |
Figure 6. The observational constraints on the variation of dark energy density as a function of redshift from WMAP and SNLS data (see [30]). The green/hatched region is excluded at 68% confidence limit, red/cross-hatched region at 95% confidence level and the blue/solid region at 99% confidence limit. The white region shows the allowed range of variation of dark energy at 68% confidence limit. |
While on the topic of observational constraints on w(t), it must be stressed that: (a) There is fair amount of tension between WMAP and SN data and one should be very careful about the priors used in these analysis. (b) There is no observational evidence for w < -1. (c) It is likely that more homogeneous, future, data sets of SN might show better agreement with WMAP results. (For more details related to these issues, see the last reference in [30].)
The observational and theoretical features described above suggests that one should consider cosmological constant as the most natural candidate for dark energy. Though it leads to well know fine tuning problems, it also has certain attractive features that need to kept in mind.
Cosmological constant is the most economical [just one number] and simplest explanation for all the observations. We stress that there is absolutely no evidence for variation of dark energy density with redshift, which is consistent with the assumption of cosmological constant.
Once we invoke the cosmological constant classical
gravity will be described by the three constants G,c and
L
-2. It is not possible to obtain a
dimensionless
quantity from these; so, within classical theory, there is no fine
tuning issue. Since
(G
/ c3)
(Lp /
L
)2
10-123, it is obvious that the cosmological
constant is telling us
something regarding quantum gravity, indicated by the
combination G
.
An acid test for any quantum gravity model
will be its ability to explain this value; needless to say, all the
currently available models - strings, loops etc. - flunk this test.
So, if dark energy is indeed cosmological constant, this will be the greatest contribution from cosmology to fundamental physics. It will be unfortunate if we miss this chance by invoking some scalar field epicycles!
In this context, it is worth stressing another peculiar feature of
cosmological constant, when it is treated as a clue to quantum gravity.
It is well known that, based on energy scales, the cosmological constant
problem is an infra red problem par excellence.
At the same time, it is a relic of a quantum gravitational effect or
principle of unknown nature. An analogy
[31]
will be helpful
to illustrate this point. Suppose one solves the Schrodinger equation
for the Helium atom for the quantum states of the two electrons
(x1,
x2). When the result is compared with observations,
one will find that only half the states - those in which
(x1,
x2) is
antisymmetric under x1
x2 interchange - are
realized in nature. But the low energy Hamiltonian for electrons in the
Helium atom has no information about
this effect! Here is low energy (IR) effect which is a relic of
relativistic quantum field theory (spin-statistics theorem) that is
totally non perturbative, in the sense that writing corrections to the
Helium atom Hamiltonian in some (1 / c) expansion will not
reproduce this result. The current value of cosmological constant could
very well be related to quantum gravity in a similar way. There must
exist a deep principle in quantum gravity which leaves its non
perturbative trace even in the low energy limit
that appears as the cosmological constant .
Let us now turn our attention to few of the many attempts to understand the cosmological constant. The choice is, of course, dictated by personal bias and is definitely a non-representative sample. A host of other approaches exist in literature, some of which can be found in [32].
One possible way of addressing this issue is to simply eliminate from
the gravitational theory those modes which couple to cosmological
constant. If, for example, we have a theory in which the source of
gravity is
( + p)
rather than (
+
3p) in Eq. (190), then
cosmological constant will not couple to gravity at all. (The non linear
coupling of matter with gravity has several subtleties; see eg.
[33].)
Unfortunately it is not possible to develop a covariant theory of
gravity using (
+ p) as the source. But we can probably gain some insight from the
following considerations. Any metric gab can be
expressed in the
form gab = f2(x)
qab such that
det q = 1 so that det g =
f4. From the action functional for gravity
![]() |
(195) |
it is obvious that the cosmological constant couples only to the
conformal factor f. So if we consider a theory of gravity in which
f4 = (-g)1/2 is kept constant and
only qab is varied, then such
a model will be oblivious of
direct coupling to cosmological constant. If the action (without the
term) is varied,
keeping det g = -1, say, then one is
lead to a unimodular theory of gravity that has the equations of
motion
Rab - (1/4)gab R =
(Tab
- (1/4)gab T) with zero trace on
both sides. Using the Bianchi identity, it is now easy to show that this
is equivalent to the usual theory with an arbitrary cosmological
constant. That is, cosmological constant arises as an undetermined
integration constant in this model
[34].
The same result arises in another, completely different approach to
gravity. In the standard approach to gravity one uses the
Einstein-Hilbert Lagrangian LEH
R which has a
formal structure LEH ~ R ~
(
g)2
+
2
g. If the surface term obtained by integrating
Lsur
2g
is ignored (or, more formally, canceled by an extrinsic curvature term)
then the Einstein's equations arise from the variation of the bulk
term Lbulk
(
g)2
which is the non-covariant
2
Lagrangian. There is, however, a remarkable relation
between Lbulk and Lsur:
![]() |
(196) |
which allows a dual description of gravity using either
Lbulk or Lsur!
It is possible to obtain the dynamics of gravity
[35]
from an approach which uses only the surface term of the Hilbert
action; we do not need the bulk term at all!. This suggests
that the true degrees of freedom of gravity
for a volume
reside in its boundary
- a point of view that
is strongly supported by the study
of horizon entropy, which shows that the degrees of freedom hidden by a
horizon scales as the area and not as the volume.
The resulting equations can be cast
in a thermodynamic form TdS = dE + PdV and the
continuum spacetime is like an elastic solid (see e.g.
[36])
with Einstein's equations
providing the macroscopic description. Interestingly, the
cosmological constant arises again in this approach as a
undetermined integration constant but closely related to the `bulk
expansion' of the solid.
While this is all very interesting, we still need an extra physical
principle to fix the value (even the sign) of cosmological constant.
One possible way of doing this is to interpret the
term in the
action as a Lagrange multiplier for the proper volume of the
spacetime. Then it is reasonable to choose the cosmological constant
such that the total proper volume of the universe is equal to a
specified number. While this will lead to a cosmological constant which
has the correct order of magnitude, it has several obvious
problems. First, the proper four volume of the universe is infinite
unless we make the spatial sections compact and restrict the range of
time integration. Second, this will lead to a dark energy density which
varies as t-2 (corresponding to w = -1/3 )
which is ruled out by observations.
Another possibility which has been attempted in the literature tries to
"cancel out" the cosmological constant by some process,
usually quantum mechanical in origin. One of the simplest ideas will be
to ask whether switching on a cosmological constant will
lead to a vacuum polarization with an effective energy momentum tensor
that will tend to cancel out the cosmological constant.
A less subtle way of doing this is to invoke another scalar field (here
we go again!) such that it can couple to
cosmological constant and reduce its effective value
[37].
Unfortunately, none of this could be made to work
properly. By and large, these approaches lead to an energy density which
is either
UV
Lp-4 (where
Lp is the Planck length) or to
IR
L
-4 (where
L
= H
-1 is the Hubble radius
associated with the cosmological constant ). The first one is too large
while the second one is too small!
8.3. Geometrical Duality in our Universe
While the above ideas do not work, it gives us a clue. A universe with two
length scales
L
and Lp will be asymptotically De Sitter
with a(t)
exp(t /
L
) at late times. There are some
curious features in such a universe which we will now describe. Given
the two length scales Lp and
L
, one can construct two energy scales
UV = 1 /
Lp4 and
IR = 1 /
L
4 in
natural units (c =
= 1). There is sufficient amount of justification
from different theoretical perspectives
to treat Lp as the zero point length of spacetime
[38],
giving a natural interpretation to
UV. The second one,
IR also has a natural
interpretation. The universe which
is asymptotically De Sitter has a horizon and associated thermodynamics
[39]
with a temperature T =
H
/ 2
and the corresponding
thermal energy density
thermal
T4
1 /
L
4 =
IR. Thus
Lp determines the highest possible energy
density in the universe while
L
determines the lowest possible energy density in this universe. As
the energy density of normal matter drops below this value, the thermal
ambience of the De Sitter phase will remain constant and provide the
irreducible `vacuum noise'. Note that the dark energy density is
the the geometric mean
DE =
(
IR
UV)1/2
between the two energy densities. If we define a
dark energy length scale LDE such that
DE = 1 /
LDE4 then
LDE = (LP
L
)1/2 is the geometric mean
of the two length scales in the universe. (Incidentally,
LDE
0.04 mm is macroscopic; it is also pretty close to the length scale
associated with a neutrino mass of 10-2 eV; another intriguing
coincidence ?!)
Using the characteristic length scale of expansion,
the Hubble radius dH
(
/ a)-1,
we can distinguish between three different phases of such a
universe. The first phase is when the universe went through a
inflationary expansion with dH =
constant; the second phase is the radiation/matter dominated phase in
which most of the standard cosmology operates and dH
increases monotonically; the third phase is that of re-inflation (or
accelerated expansion) governed by the cosmological constant in which
dH is again
a constant. The first and last phases are time translation invariant;
that is, t
t + constant is an (approximate) invariance for the
universe in these two phases. The universe satisfies the perfect
cosmological principle and is in steady state during these phases!
In fact, one can easily imagine a scenario in which the two De Sitter
phases (first and last) are of arbitrarily long duration
[40].
If
0.7,
DM
0.3
the final De Sitter phase does last forever; as regards the
inflationary phase, nothing prevents it from lasting for arbitrarily
long duration. Viewed from this perspective, the in between phase -
in which most of the `interesting' cosmological phenomena occur - is
of negligible measure in the span of time. It merely connects two
steady state phases of the universe.
While the two De Sitter phases can last forever in principle, there is a
natural cut off length scale in both of them
which makes the region of physical relevance to be finite
[40].
Let us first discuss the case of re-inflation in the
late universe.
As the universe grows exponentially in the phase 3, the wavelength of
CMBR photons are being redshifted rapidly. When the temperature of the
CMBR radiation drops below the De Sitter temperature (which happens when
the wavelength of the typical CMBR photon is stretched to the
L.)
the universe will be essentially dominated by the vacuum thermal noise
[39]
due to the horizon in the De Sitter phase.
This happens when the expansion factor is a =
af determined by the
equation T0 (a0 /
af) = (1/2
L
). Let a =
a
be the epoch at which
cosmological constant started dominating over matter, so that
(a
/ a0)3
= (
DM
/
).
Then we find that the dynamic range of the phase 3 is
![]() |
(197) |
Interestingly enough, one can also impose a similar bound on the
physically relevant duration of inflation. We know that the quantum
fluctuations, generated during this inflationary phase, could act as
seeds of structure formation in the universe. Consider a perturbation
at some given wavelength scale which is stretched with the expansion of
the universe as
a(t).
During the inflationary phase, the Hubble radius remains constant while
the wavelength increases, so that the perturbation will `exit' the
Hubble radius at some time. In the radiation dominated phase, the Hubble
radius dH
t
a2
grows faster than the wavelength
a(t). Hence, normally, the perturbation will `re-enter'
the Hubble radius at some time.
If there was no re-inflation, all wavelengths will re-enter the
Hubble radius sooner or later.
But if the universe undergoes re-inflation, then the Hubble radius
`flattens out' at late times and some of the perturbations will
never reenter the Hubble radius ! If we use the criterion that we
need the perturbation to reenter the Hubble radius, we get a natural
bound on the duration of inflation which is of direct astrophysical
relevance. Consider a perturbation which leaves the Hubble radius
(Hin-1) during the inflationary
epoch at a = ai. It will grow
to the size Hin-1(a /
ai) at a later epoch.
We want to determine ai such that this length scale
grows to
L
just when the dark energy starts dominating over
matter; that is at the epoch a =
a
=
a0(
DM /
)1/3. This gives
Hin-1(a
/ ai) =
L
so that
ai = (Hin-1
/ L
)(
DM /
)1/3 a0. On
the other hand, the inflation ends at
a = aend where aend /
a0 = T0 / Treheat
where Treheat is the temperature to which the
universe has been reheated at the end of inflation. Using these two
results we can determine the dynamic range of this phase 1 to be
![]() |
(198) |
where we have used the fact that, for a GUTs scale inflation with
EGUT = 1014 GeV, Treheat
= EGUT,
in
= EGUT4
we have 2
Hin-1 Treheat
= (3
/
2)1/2(Ep / EGUT)
105.
If we consider a quantum gravitational, Planck scale, inflation with
2
Hin-1 Treheat =
(1), the ranges in
Eq. (197) and Eq. (198) are approximately equal.
This fact is definitely telling us something regarding the duality between Planck scale and Hubble scale or between the infrared and ultraviolet limits of the theory. The mystery is compounded by the fact the asymptotic De Sitter phase has an observer dependent horizon and related thermal properties [39]. Recently, it has been shown - in a series of papers, see ref. [35] - that it is possible to obtain classical relativity from purely thermodynamic considerations. It is difficult to imagine that these features are unconnected and accidental; at the same time, it is difficult to prove a definite connection between these ideas and the cosmological constant.
8.4. Gravity as detector of the vacuum energy
Finally, we will describe an idea which does lead to the
correct value of cosmological constant.
The conventional discussion of the relation between cosmological
constant and vacuum energy density is based on
evaluating the zero point energy of quantum fields with an ultraviolet
cutoff and using the result as a source of gravity.
Any reasonable cutoff will lead to a vacuum energy density
vac
which is unacceptably high. This argument, however, is too simplistic
since the zero point energy - obtained by summing over the
(1/2)
k -
has no observable consequence in any other
phenomena and can be subtracted out by redefining the Hamiltonian. The
observed non trivial features of the vacuum state of QED, for example,
arise from the fluctuations (or modifications) of this vacuum
energy rather than the vacuum energy itself.
This was, in fact, known fairly early in the history of cosmological
constant problem and is stressed by Zeldovich
[41]
who explicitly calculated one possible contribution to fluctuations
after subtracting away the mean value. This
suggests that we should consider the fluctuations in the vacuum energy
density in addressing the cosmological constant problem.
If the vacuum probed by the gravity can readjust to take away the bulk
energy density
UV
Lp-4, quantum fluctuations can generate
the observed value
DE.
One of the simplest models
[42]
which achieves this uses the fact that, in the
semi-classical limit, the wave function describing the universe of
proper four-volume
will
vary as
exp(-iA0)
exp[ -i(
eff
/
Lp2)]. If we treat
(
/
Lp2,
) as conjugate variables
then uncertainty principle suggests
Lp2 /
. If
the four volume is built out of Planck scale substructures, giving
=
NLp4, then the Poisson fluctuations will
lead to
1/2
Lp2 giving
= Lp2 /
1 /
1/2
H02 . (This idea can be a more
quantitative; see
[42]).
Similar viewpoint arises, more rigorously, when we study the question of
detecting the energy
density using gravitational field as a probe.
Recall that an Unruh-DeWitt detector with a local coupling
LI =
M()
[x(
)] to the field
actually responds to
<0|
(x)
(y)|0> rather than to the field itself
[43].
Similarly, one can use the gravitational
field as a natural "detector" of energy momentum tensor
Tab with the standard coupling L =
hab
Tab. Such a model was analysed in detail in ref.
[44]
and it was shown that the
gravitational field responds to the two point function 〈
0|Tab(x) Tcd(y)|0〉 .
In fact, it is essentially this
fluctuations in the energy density which is computed in the inflationary
models (see Eq. (170)) as the seed source for gravitational
field, as stressed in ref.
[11].
All these suggest treating the energy fluctuations as
the physical quantity "detected" by gravity, when
one needs to incorporate quantum effects.
If the cosmological constant arises due to the energy density of the
vacuum, then one needs to understand the structure of the quantum vacuum
at cosmological scales. Quantum theory, especially the paradigm of
renormalization group has taught us that the energy density - and even
the concept of the vacuum state - depends on the scale at which it is
probed. The vacuum state which we use to study the lattice vibrations in
a solid, say, is not the same as vacuum state of the QED.
In fact, it seems inevitable that in a universe with two length
scale L, Lp, the vacuum
fluctuations will contribute an energy density of the correct order of
magnitude
DE =
(
IR
UV)1/2.
The hierarchy of energy scales in such a universe, as detected
by the gravitational field has
[40,
45]
the pattern
![]() |
(199) |
The first term is the bulk energy density which needs to be renormalized away (by a process which we do not understand at present); the third term is just the thermal energy density of the De Sitter vacuum state; what is interesting is that quantum fluctuations in the matter fields inevitably generate the second term.
The key new ingredient arises from the fact that the properties of the
vacuum state depends on the scale at which it is probed and it is not
appropriate to ask questions without specifying this scale.
If the spacetime has a cosmological horizon which blocks information,
the natural scale is provided by the size of the horizon,
L,
and we should use observables defined within the accessible region.
The operator H(<
L
), corresponding to the total energy inside
a region bounded by a cosmological horizon, will exhibit fluctuations
E since
vacuum state is not an eigenstate of
this operator. The corresponding fluctuations in the energy
density,
(
E)
/L
3 = f(Lp,
L
)
will now depend on both the ultraviolet cutoff Lp as
well as
L
. To obtain
vac
E /
L
3 which scales as
(Lp L
)-2 we need to have
(
E)2
Lp-4 L
2; that is,
the square of the energy fluctuations should scale as the surface area
of the bounding surface which is provided by the cosmic
horizon. Remarkably enough, a rigorous calculation
[45]
of the dispersion in the energy shows that
for L
>> Lp, the final result
indeed has the scaling
![]() |
(200) |
where the constant c1 depends on the manner in which
ultra violet cutoff is imposed.
Similar calculations have been done (with a completely different
motivation, in the context of entanglement entropy)
by several people and it is known that the area scaling found in
Eq. (200), proportional to
L2, is a generic feature
[46].
For a simple exponential UV-cutoff, c1 =
(1/30
2) but cannot be
computed reliably without knowing the full theory.
We thus find that the fluctuations in the energy density of the vacuum
in a sphere of radius
L
is given by
![]() |
(201) |
The numerical coefficient will depend on c1 as well as
the precise nature of infrared cutoff
radius (like whether it is L or
L
/ 2
etc.). It
would be pretentious to cook up the factors
to obtain the observed value for dark energy density.
But it is a fact of life that a fluctuation of magnitude
vac
H
2 / G will exist in the
energy density inside a sphere of radius
H
-1 if Planck
length is the UV cut off. One cannot get away from it.
On the other hand, observations suggest that there is a
vac of
similar magnitude in the universe. It seems
natural to identify the two, after subtracting out the mean value by
hand. Our approach explains why there is a surviving
cosmological constant which satisfies
DE
= (
IR
UV)1/2
which - in our opinion - is the problem.
Acknowledgement
I thank J. Alcaniz for his friendship and warm hospitality during the X Special Courses at Observatorio Nacional, Rio de Janeiro, Brazil and for persuading me to write up my lecture notes. I am grateful to Gaurang Mahajan for help in generating the figures.