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Journal of Gravity

Volume 2013 (2013), Article ID 682451, 20 pages

http://dx.doi.org/10.1155/2013/682451

## A Cosmological Model Based on a Quadratic Equation of State Unifying Vacuum Energy, Radiation, and Dark Energy

Laboratoire de Physique Théorique, IRSAMC, CNRS, UPS, Université de Toulouse, 31062 Toulouse, France

Received 26 March 2013; Accepted 17 May 2013

Academic Editor: Kazuharu Bamba

Copyright © 2013 Pierre-Henri Chavanis. This is an open access article distributed under the Creative Commons Attribution License, which permits unrestricted use, distribution, and reproduction in any medium, provided the original work is properly cited.

#### Abstract

We consider a cosmological model based on a quadratic equation of state (where is the Planck density and is the cosmological density) “unifying” vacuum energy, radiation, and dark energy. For , it reduces to leading to a phase of early accelerated expansion (early inflation) with a constant density equal to the Planck density g/m^{3} (vacuum energy). For , we recover the equation of state of radiation . For , we get leading to a phase of late accelerated expansion (late inflation) with a constant density equal to the cosmological density g/m^{3} (dark energy). The temperature is determined by a generalized Stefan-Boltzmann law. We show a nice “symmetry” between the early universe (vacuum energy + radiation) and the late universe (radiation + dark energy). In our model, they are described by two polytropic equations of state with index and respectively. Furthermore, the Planck density in the early universe plays a role similar to that of the cosmological density in the late universe. They represent fundamental upper and lower density bounds differing by 122 orders of magnitude. We add the contribution of baryonic matter and dark matter considered as independent species and obtain a simple cosmological model describing the whole evolution of the universe. We study the evolution of the scale factor, density, and temperature. This model gives the same results as the standard CDM model for , where is the Planck time and completes it by incorporating the phase of early inflation in a natural manner. Furthermore, this model does not present any singularity at and exists eternally in the past (although it may be incorrect to extrapolate the solution to the infinite past). Our study suggests that vacuum energy, radiation, and dark energy may be the manifestation of a unique form of “generalized radiation.” By contrast, the baryonic and dark matter components of the universe are treated as different species. This is at variance with usual models (quintessence, Chaplygin gas, ...) trying to unify dark matter and dark energy.

#### 1. Introduction

The evolution of the universe may be divided into four main periods [1]. In the vacuum energy era (Planck era), the universe undergoes a phase of early inflation that brings it from the Planck size m to an almost “macroscopic” size m in a tiniest fraction of a second [2–5]. The universe then enters in the radiation era and, when the temperature cools down below approximately 10^{3} K, in the matter era [6]. Finally, in the dark energy era (de Sitter era), the universe undergoes a phase of late inflation [7]. The early inflation is necessary to solve notorious difficulties such as the singularity problem, the flatness problem, and the horizon problem [2–5]. The late inflation is necessary to account for the observed accelerating expansion of the universe [8–11]. At present, the universe is composed of approximately 5% baryonic matter, 20% dark matter, and 75% dark energy [1]. Despite the success of the standard model, the nature of vacuum energy, dark matter, and dark energy remains very mysterious and leads to many speculations.

The phase of inflation in the early universe is usually described by some hypothetical scalar field with its origin in quantum fluctuations of vacuum [2–5]. This leads to an equation of state , implying a constant energy density, called the vacuum energy. This energy density is usually identified with the Planck density . As a result of the vacuum energy, the universe expands exponentially rapidly on a timescale of the order of the Planck time . This phase of early inflation is followed by the radiation era described by an equation of state . In the vacuum energy era, the density has a constant value , while it decreases as during the radiation era.

The phase of acceleration in the late universe is usually ascribed to the cosmological constant which is equivalent to a constant energy density called the dark energy. This acceleration can be modeled by an equation of state , implying a constant energy density identified with the cosmological density . As a result of the dark energy, the universe expands exponentially rapidly on a timescale of the order of the cosmological time (de Sitter solution). This leads to a phase of late inflation. Inspired by the analogy with the early inflation, some authors have represented the dark energy by a scalar field called quintessence [12–24]. As an alternative to the quintessence, other authors have proposed to model the acceleration of the universe by an exotic fluid with an equation of state of the form with and called the Chaplygin gas when and the generalized Chaplygin gas otherwise [25–32]. This equation of state may be rewritten as with in order to emphasize its analogy with the polytropic equation of state [33–36]. Indeed, it may be viewed as the sum of a linear equation of state describing dust matter () or radiation () and a polytropic equation of state . We shall call (1) a generalized polytropic equation of state. At late times, this equation of state with leads to a constant energy density implying an exponential growth of the scale factor that is similar to the effect of the cosmological constant. At earlier times, and for , this equation of state returns the results of the dust matter model (). Therefore, this hydrodynamic model provides a unification of matter () and dark energy () in the late universe [25–36].

In [34–36], we remarked that a constant pressure has the same virtues as the Chaplygin gas model. (Actually, this is a particular case of the generalized Chaplygin gas model corresponding to and .) This equation of state may be viewed as a generalized polytropic equation of state (1) with polytropic index , polytropic constant , and linear coefficient . At late times, this equation of state leads to a constant energy density equal to implying an exponential growth of the scale factor (late inflation). At earlier times, this equation of state returns the results of the dust matter model (). Therefore, as any Chaplygin gas model with and , this equation of state provides a unification of matter () and dark energy () in the late universe. Moreover, this equation of state gives the *same* results as the standard CDM model where matter and dark energy are treated as independent species. This is interesting since the CDM model is consistent with observations and provides a good description of the late universe. (By contrast, the Chaplygin gas model is not consistent with observations unless is close to zero [37]. Therefore, observations tend to select the constant pressure equation of state (, ) among the whole family of Chaplygin gas models of the form .) In [34–36], we also proposed to describe the transition between the vacuum energy era and the radiation era in the early universe by a unique equation of state of the form . This equation of state may be viewed as a generalized polytropic equation of state (1) with polytropic index , polytropic constant , and linear coefficient . At early times, this equation of state leads to a constant energy density equal to implying an exponential growth of the scale factor (early inflation). At later times, and for , this equation of state returns the results of the radiation model (). Therefore, this hydrodynamic model provides a unification of vacuum energy () and radiation () in the early universe. This equation of state may be viewed as an extension of the Chaplygin (or polytropic) gas model to the early universe. In this approach, the late universe is described by negative polytropic indices (e.g., ) and the early universe by positive polytropic indices (e.g., ) [34–36]. Indeed, for , the polytropic equation of state dominates in the late universe where the density is low, and for , the polytropic equation of state dominates in the early universe where the density is high.

In the viewpoint just exposed, one tries to “unify” vacuum energy + radiation on the one hand and dust matter + dark energy on the other hand. (It is oftentimes argued that the dark energy (cosmological constant) corresponds to the vacuum energy. This leads to the so-called *cosmological problem* [38] since the cosmological density and the Planck density differ by about 122 orders of magnitude. We think that it is a mistake to identify the dark energy with the vacuum energy. In this paper, we shall regard the vacuum energy and the dark energy as two distinct entities. We shall call vacuum energy the energy associated with the Planck density and dark energy the energy associated with the cosmological density. The vacuum energy is responsible for the inflation in the early universe and the dark energy for the inflation in the late universe. In this viewpoint, the vacuum energy is due to quantum mechanics and the dark energy is an effect of general relativity. The cosmological constant is interpreted as a fundamental constant of nature applying to the cosmophysics in the same way the Planck constant applies to the microphysics.) In this paper, we adopt a different point of view. We propose that vacuum energy, radiation, and dark energy may be the manifestation of a unique form of “generalized radiation”. We propose to describe it by a *quadratic* equation of state
involving the Planck density and the cosmological density. In the early universe, we recover the equation of state unifying vacuum energy and radiation. In the late universe, we obtain an equation of state unifying radiation and dark energy. This equation of state may be viewed as a generalized polytropic equation of state (1) with polytropic index , polytropic constant , and linear coefficient . The quadratic equation of state (2) “unifies” vacuum energy, radiation, and dark energy in a natural and simple manner. It also allows us to obtain a generalization of the Stefan-Boltzmann law for the evolution of the temperature [see (69)] that is valid during the whole evolution of the universe. In this approach, baryonic matter and dark matter are treated as different species. This is at variance with the usual scenario where vacuum energy + radiation on the one hand and dark matter + dark energy on the other hand are treated as different species.

The paper is organized as follows. In Section 2, we recall the basic equations of cosmology that will be needed in our study. In Section 3, we describe the transition between the vacuum energy era and the radiation era in the early universe. In Section 4, we describe the transition between the radiation era and the dark energy era in the late universe. In Section 5, we introduce the general model where vacuum energy + radiation + dark energy are described by a quadratic equation of state while baryonic matter and dark matter are considered as different species. We determine the general equations giving the evolution of the scale factor, density, and temperature (Section 5.1). We consider particular limits of these equations in the early universe (Section 5.2) and in the late universe (Sections 5.3 and 5.4). We connect these two limits and describe the whole evolution of the universe (Section 5.5). This model gives the same results as the standard CDM model for and completes it by incorporating the early inflation in a natural manner [see (100)]. It also reveals a nice “symmetry” between the early universe (vacuum energy + radiation) and the late universe (radiation + dark energy). These two phases are described by polytropic equations of state with index and , respectively. Furthermore, the cosmological density in the late universe plays a role similar to the Planck density in the early universe. They represent fundamental upper and lower density bounds differing by orders of magnitude.

#### 2. Basic Equations of Cosmology

In a space with uniform curvature, the line element is given by the Friedmann-Lemaître-Roberston-Walker (FLRW) metric where represents the radius of curvature of the -dimensional space, or the scale factor. (In our conventions, and are dimensionless, while the scale factor has the dimension of a length. It is defined such that (where is the speed of light and is the Hubble constant) at the present epoch. This corresponds to the cosmological horizon. The usual (dimensionless) scale factor is . It satisfies at the present epoch.) By an abuse of language, we shall sometimes call it the “radius of the universe.” On the other hand, determines the curvature of space. The universe may be closed (), flat (), or open ().

If the universe is isotropic and homogeneous at all points in conformity with the line element (3) and contains a uniform perfect fluid of energy density and isotropic pressure , the energy-momentum tensor is The Einstein equations relate the geometrical structure of the spacetime () to the material content of the universe (). For the sake of generality, we have accounted for a possibly nonzero cosmological constant . Given Equations (3) and (4), these equations reduce to where dots denote differentiation with respect to time. These are the well-known cosmological equations describing a non-static universe first derived by Friedmann [6].

The Friedmann equations are usually written in the form where we have introduced the Hubble parameter . Among these three equations, only two are independent. The first equation, which can be viewed as an “equation of continuity”, can be directly derived from the conservation of the energy momentum tensor , which results from the Bianchi identities. For a given barotropic equation of state , it determines the relation between the density and the scale factor. Then, the temporal evolution of the scale factor is given by (9).

Equivalent expressions of the “equation of continuity” are Introducing the volume and the energy , (11) becomes . This can be viewed as the first principle of thermodynamics for an adiabatic evolution of the universe .

From the first principle of thermodynamics in the form , one can derive the thermodynamical equation [6] For a given barotropic equation of state , this equation can be integrated to obtain the relation between the temperature and the density. Combining (12) with (13), we get [6] that is the entropy of the universe in a volume . This confirms that the Friedmann equations (7)–(9) imply the conservation of the entropy.

In this paper, we consider a flat universe () in agreement with the observations of the cosmic microwave background (CMB) [1]. On the other hand, we set because the effect of the cosmological constant will be taken into account in the equation of state. The Friedmann equations then reduce to The deceleration parameter is defined by The universe is decelerating when and accelerating when . Introducing the equation of state parameter and using the Friedmann equations (16) and (17), we obtain for a flat universe We see from (19) that the universe is decelerating if (strong energy condition) and accelerating if . On the other hand, according to (15), the density decreases with the scale factor if (null dominant energy condition) and increases with the scale factor if .

#### 3. The Early Universe

In [34–36], we have proposed to describe the transition between the vacuum energy era and the radiation era in the early universe by a single equation of state of the form where is the Planck density. This equation of state corresponds to a generalized polytropic equation of state (1) with , , and . For , we recover the equation of state of the pure radiation . For , we get corresponding to the vacuum energy.

The continuity equation (15) may be integrated into
where is a constant of integration. This characteristic scale marks the transition between the vacuum energy era and the radiation era. The equation of state (20) interpolates smoothly between the vacuum energy era ( and ) and the radiation era ( and ). It provides therefore a “unified” description of vacuum energy and radiation in the early universe. This amounts to summing the *inverse* of the densities of these two phases. Indeed, (21) may be rewritten as
At we have so that .

The equation of state parameter and the deceleration parameter are given by The velocity of sound is given by As the universe expands from to , the density decreases from to , the equation of state parameter increases from to , the deceleration parameter increases from to , and the ratio increases from to .

##### 3.1. The Vacuum Energy Era: Early Inflation

When , the density tends to a maximum value
and the pressure tends to . The Planck density (vacuum energy) represents a fundamental upper bound for the density. A constant value of the density gives rise to a phase of early inflation. From the Friedmann equation (17), we find that the Hubble parameter is constant where we have introduced the Planck time . Numerically, . Therefore, the scale factor increases exponentially rapidly with time as
The timescale of the exponential growth is the Planck time . We have defined the “original” time such that is equal to the Planck length . Mathematically speaking, the universe exists at any time in the past ( and for ), so there is no primordial singularity. However, when , we cannot ignore quantum fluctuations associated with the spacetime. In that case, we cannot use the classical Einstein equations anymore and a theory of quantum gravity is required. It is not known whether quantum gravity will remove, or not, the primordial singularity. Therefore, we cannot extrapolate the solution (26) to the infinite past. However, this solution may provide a *semiclassical* description of the phase of early inflation when .

##### 3.2. The Radiation Era

When , we recover the equation corresponding to the pure radiation described by an equation of state . The conservation of implies that where is the present density of radiation and the present distance of cosmological horizon determined by the Hubble constant (the Hubble time is ). Writing where is the present density of the universe given by (17) and is the present fraction of radiation in the universe [1], we obtain This scale is intermediate between the Planck length and the present size of the universe ( and ). It gives the typical size of the universe at the end of the inflationary phase (or at the beginning of the radiation era). When , the Friedmann equation (17) yields We then have During the radiation era, the density decreases algebraically as the universe expands.

##### 3.3. The General Solution

The general solution of the Friedmann equation (17) with (21) is where is a constant of integration. It is determined such that at . Setting , we get Numerically, and . For , we have the exact asymptotic result with . Due to the smallness of , a very good approximation of and is given by and . With this approximation, (33) returns (26). Using (23), the universe is accelerating when (i.e., ) and decelerating when (i.e., ) where and . The time at which the universe starts decelerating is given by . This corresponds to the time at which the curve presents an inflexion point. It turns out that this inflexion point coincides with so it also marks the end of the inflation (). Numerically, , , and . At , and . Numerically, and .

##### 3.4. The Temperature and the Pressure

The thermodynamical equation (13) with the equation of state (20) may be integrated into where is a constant of integration. Using (21), we get Equation (34) may be seen as a generalized Stefan-Boltzmann law in the early universe. In the radiation era , it reduces to the ordinary Stefan-Boltzmann law This allows us to identify the constant of integration with the Planck temperature . On the other hand, in the inflationary phase , the temperature is related to the density by According to (35), for and for .

The temperature starts from at , increases exponentially rapidly during the inflation as reaches a maximum value and decreases algebraically as during the radiation era. Using (34) and (35), we find that the point corresponding to the maximum temperature is , , . It is reached at a time . Numerically, , , , and . At , . Numerically, . At , . Numerically, .

The pressure is given by (20). Using (21), we get The pressure starts from at , remains approximately constant during the inflation, becomes positive, reaches a maximum value , and decreases algebraically during the radiation era. At , . Numerically, . The point at which the pressure vanishes is , , . This corresponds to a time . Numerically, , , , and . On the other hand, the pressure reaches its maximum at the same point as the one at which the temperature reaches its maximum. Numerically, . At , we have . Numerically, .

Using (20), (21), and (35), we find that the entropy (14) is given by and we check that it is a constant. Numerically, .

##### 3.5. The Evolution of the Early Universe

In our model, the universe starts at with a vanishing radius , a finite density , a finite pressure , and a vanishing temperature . The universe exists at any time in the past and does not present any singularity. For , the radius of the universe is less than the Planck length . In the Planck era, quantum gravity should be taken into account and our semiclassical approach is probably not valid in the infinite past. At , the radius of the universe is equal to the Planck length . The corresponding density, temperature, and pressure are , , and . We note that quantum mechanics regularizes the finite time singularity present in the standard Big Bang theory. This is similar to finite size effects in second-order phase transitions (see Section 3.6). The Big Bang theory is recovered for . We also note that the universe is very cold at , contrary to what is predicted by the Big Bang theory (a naive extrapolation of the law leads to ). The universe first undergoes a phase of inflation during which its radius and temperature increase exponentially rapidly, while its density and pressure remain approximately constant. The inflation “starts” at and ends at . During this very short lapse of time, the radius of the universe grows from to , and the temperature grows from to . By contrast, the density and the pressure do not change significatively: they go from and to and . The pressure passes from negative values to positive values at corresponding to , , and . After the inflation, the universe enters in the radiation era and, from that point, we recover the standard model [6]. The radius increases algebraically as , while the density and the temperature decrease algebraically as and . The temperature and the pressure achieve their maximum values and at . At that moment, the density is and the radius . During the inflation, the universe is accelerating and during the radiation era it is decelerating. The transition (marked by an inflexion point) takes place at a time coinciding with the end of the inflation (). The evolution of the scale factor, density, and temperature as a function of time are represented in Figures 1, 2, 3, 4, 5, 6, and 7 in logarithmic and linear scales.

We note that the inflationary process described previously is relatively different from the usual inflationary scenario [2–5]. In standard inflation, the universe is radiation dominated up to but expands exponentially by a factor in the interval with . For , the evolution is again radiation dominated. At , the temperature is about (this corresponds to the epoch at which most “grand unified theories” have a significant influence on the evolution of the universe). During the exponential inflation, the temperature drops drastically; however, the matter is expected to be reheated to the initial temperature of by various high energy processes [39]. In the inflationary process described previously, the evolution of the temperature, given by a generalized Stefan-Boltzmann law, is different. It is initially very low and increases during the inflation.

##### 3.6. Analogy with Phase Transitions

The standard Big Bang theory is a classical theory in which quantum effects are neglected. In that case, it exhibits a finite time singularity: the radius of the universe is equal to zero at , while its density and its temperature are infinite. For , the solution is not defined and we may take . For , the radius of the universe increases as . This is similar to a second-order phase transition if we view the time as the control parameter (e.g., the temperature ) and the scale factor as the order parameter (e.g., the magnetization ). The exponent is the same as in mean field theories of second-order phase transitions (i.e., ), but this is essentially a coincidence.

When quantum mechanics effects are taken into account, as in our simple semiclassical model, the singularity at disappears and the curves , and are regularized. In particular, we find that at , instead of , due to the finite value of . This is similar to the regularization due to finite size effects (e.g., the system size or the number of particles ) in ordinary phase transitions. In this sense, the classical limit is similar to the thermodynamic limit ( or ) in ordinary phase transitions.

To study the convergence toward the classical Big Bang solution when , it is convenient to use a proper normalization of the parameters. We call the true value of the Planck constant (in this section, the subscript refers to true values of the parameters). Then, we express time in terms of and lengths in terms of . We introduce and . Using and , we obtain where . Therefore where . Using these relations and (31), we get where is defined by (32). This equation describes a phase transition between the vacuum energy era and the radiation era. The normalized Planck constant plays the role of finite size effects. Finite size scalings are explicit in (44). For , we recover the nonsingular model of Section 3.3. For , we recover the singular radiation model of Section 3.2 (Big Bang). The convergence towards this singular solution as is shown in Figure 8.

##### 3.7. Scalar Field Theory

The phase of inflation in the very early universe is usually described by a scalar field [5]. The ordinary scalar field minimally coupled to gravity evolves according to the equation where is the potential of the scalar field. The scalar field tends to run down the potential towards lower energies. The density and the pressure of the universe are related to the scalar field by Using standard technics [7], we find that the potential corresponding to the equation of state (20) is [34–36] In the vacuum energy era (), using (26), we get In the radiation era (), using (29), we get More generally, using (31), the evolution of the scalar field in the early universe is given by These results are illustrated in Figures 9 and 10.

#### 4. The Radiation and the Dark Energy in the Late Universe

We propose to describe the transition between the radiation era and the dark energy era in the late universe by a single equation of state of the form where is the cosmological density. This equation of state corresponds to a generalized polytropic equation of state (1) with , , and . It may be viewed as the “symmetric” version of the equation of state (20) in the early universe. For , we recover the equation of state of the pure radiation . For , we get corresponding to the dark energy.

The continuity equation (15) may be integrated into where is a constant of integration. This characteristic scale marks the transition between the radiation era and the dark energy era. The equation of state (51) interpolates smoothly between the radiation era ( and ) and the dark energy era ( and ). It provides therefore a “unified” description of radiation and dark energy in the late universe. This amounts to summing the density of these two phases. Indeed, (52) may be rewritten as At we have so that .

##### 4.1. The Dark Energy Era: Late Inflation

When , the density tends to a minimum value and the pressure tends to . The cosmological density (dark energy) represents a fundamental lower bound for the density. A constant value of the density gives rise to a phase of late inflation. It is convenient to define a cosmological time and a cosmological length . These are the counterparts of the Planck scales for the late universe. From the Friedmann equation (17), we find that the Hubble parameter is constant . Numerically, . Therefore, the scale factor increases exponentially rapidly with time as This exponential growth corresponds to the de Sitter solution [1]. The timescale of the exponential growth is the cosmological time . This solution exists at any time in the future ( and for ), so there is no future singularity. This is not the case of all cosmological models. In a “phantom universe” [40–57], violating the null dominant energy condition (), the density increases as the universe expands. This may lead to a future singularity called Big Rip (the density becomes infinite in a finite time). The possibility that we live in a phantom universe is not ruled out by observations.

##### 4.2. The Radiation Era

When , we recover the equation corresponding to the pure radiation described by an equation of state . The conservation of implies that . Using and with (dark energy) and (radiation) [1], we obtain hence This can be rewritten as where we have used . This scale gives the typical size of the universe at the transition between the radiation era and the dark energy era. When , the Friedmann equation (17) yields We then have

##### 4.3. The Temperature and the Pressure

The thermodynamical equation (13) with the equation of state (51) may be integrated into
where is a constant of integration. Using (52), we get
Equation (60) may be seen as a generalized Stefan-Boltzmann law in the late universe. In the radiation era , it reduces to the ordinary Stefan-Boltzmann law (36). This allows us to identify the constant of integration with the Planck temperature . We note that both in the radiation *and* dark energy eras. This suggests that radiation and dark energy belong to the same “family”. In the radiation era the temperature decreases algebraically as , and in the dark energy era, it decreases exponentially rapidly as .

The pressure is given by (51). Using (52), we get It decreases algebraically during the radiation era and tends to a constant negative value for . Numerically, . The point at which the pressure vanishes is , , . Numerically, , , .

#### 5. The General Model

##### 5.1. The Quadratic Equation of State

We propose to describe the vacuum energy, radiation, and dark energy by a unique equation of state For , (vacuum energy); for , (radiation); for , (dark energy). This quadratic equation of state combines the properties of the equation of state (20) valid in the early universe and of the equation of state (51) valid in the late universe. A nice feature of this equation of state is that both the Planck density (vacuum energy) and the cosmological density (dark energy) explicitly appear. Therefore, this equation of state reproduces both the early inflation and the late inflation. On the other hand, this description suggests that vacuum energy, radiation, and dark energy may be the manifestation of a unique form of “generalized radiation”. This is, however, just a speculation before a justification of this equation of state has been given from first principles. In this paper, we limit ourselves to studying the properties of this equation of state.

Using the equation of continuity (15), we obtain where (see Section 3). To obtain this expression, we have used the fact that so that . When , (vacuum energy); when , (radiation); when , (dark energy). In the early universe, the contribution of dark energy is negligible and we recover (21). In the late universe, the contribution of vacuum energy is negligible and we recover (52) with .

We shall view the “generalized radiation”, the baryonic matter, and the dark matter as different species. In the general case, these species could interact. However, in this paper, we assume, for simplicity, that they do not. The generalized radiation (vacuum energy + radiation + dark energy) is described by the quadratic equation of state (63). Using (see Section 3), we can rewrite (64) in the form where . The baryonic matter and the dark matter are described as pressureless fluids with and . The equation of continuity (15) for each species leads to the relations

The total density of generalized radiation, baryonic matter, and dark energy can be written as Substituting this relation in the Friedmann equation (17) and writing , , , and , we obtain From observations, , , , , and [1].

The thermodynamic equation (13) with the equation of state (63) may be integrated into Using (64), we get Equation (70) may also be written as with . Numerically, . Since , can be identified with the present temperature of radiation. Indeed, after the vacuum energy era, (71) reduces to

Finally, using (63) and (64), the pressure of generalized radiation can be written in very good approximation as

##### 5.2. The Early Universe

In the early universe, we can neglect the contribution of baryonic matter, dark matter, and dark energy in the density. We only consider the contribution of vacuum energy and radiation. Equation (67) then reduces to The Friedmann equation (68) becomes It has the analytical solution given by (31). On the other hand, the temperature of the generalized radiation in the early universe is given by (34) and (35). The transition between the vacuum energy era and the radiation era corresponds to . This yields , , . Numerically, , , . This takes place at a time . This time also corresponds to the inflexion point of the curve . The universe is accelerating for and decelerating for .

In the vacuum energy era (), the density is constant: Therefore, the Hubble parameter is also a constant . In that case, the scale factor and the temperature increase exponentially rapidly according to (26) and (38).

In the radiation era (), we have

The Friedmann equation (75) reduces to and yields

##### 5.3. The Late Universe

In the late universe, we can neglect the contribution of vacuum energy and radiation in the density. We only consider the contribution of baryonic matter, dark matter, and dark energy. Equation (67) then reduces to where . This coincides with the CDM model (see Section 5.5). The Friedmann equation (68) becomes with . It has the analytical solution The density evolves as On the other hand, the temperature of the generalized radiation in the late universe is given by Using the expression (82) of the scale factor, we obtain Finally, the pressure of the generalized radiation in the late universe may be written as

Setting in (82), we find the age of the universe Numerically, . It is rather fortunate that the age of the universe almost coincides with the Hubble time .

The universe starts accelerating when [see (16) with and ]. Using and with (matter) and (dark energy), we obtain and with At that point and . Numerically, , , , and . The universe is decelerating for and accelerating for .

The transition between the matter era and the dark energy era corresponds to . This leads to and with At that point and . Numerically, , , , and . We also note that . Of course, the value coincides with the one obtained in [34–36] using a different viewpoint.

In the dark energy era (), the density is constant: Therefore, the Hubble parameter is also a constant . Using (82), we find that the scale factor increases exponentially rapidly as This corresponds to the de Sitter solution. According to (85), the temperature decreases exponentially rapidly as In the matter era (), we have The Friedmann equation (81) reduces to and we have This corresponds to the Einstein-de Sitter (EdS) universe.

##### 5.4. The Evolution of the Late Universe

In the matter era (), the radius increases algebraically as , while the density decreases algebraically as (EdS). When , the universe enters in the dark energy era. It undergoes a late inflation (de Sitter) during which its radius increases exponentially rapidly, while its density remains constant and equal to the cosmological density . The transition takes place at , , , and . In the matter era, the universe is decelerating, and in the dark energy era, it is accelerating. The time at which the universe starts accelerating is , corresponding to a radius , a density and a temperature . In the late universe, the temperature of generalized radiation always scales as . It decreases algebraically rapidly () in the matter era and exponentially rapidly in the dark energy era. The evolution of the scale factor, density, temperature, and pressure of radiation as a function of time is represented in Figures 11, 12, 13, 14, 15, and 16 in logarithmic and linear scales.