Advertisement

Study of parametrized dark energy models with a general non-canonical scalar field

Open Access
Regular Article - Theoretical Physics

Abstract

In this paper, we consider various dark energy models in the framework of a non-canonical scalar field with a Lagrangian density of the form \(\mathcal{L}(\phi , X)=f(\phi )X{\left( \frac{X}{M^{4}_{Pl}}\right) }^{\alpha -1} - V(\phi )\), which provides the standard canonical scalar field model for \(\alpha =1\) and \(f(\phi )=1\). In this particular non-canonical scalar field model, we carry out the analysis for \(\alpha =2\). We then obtain cosmological solutions for constant as well as variable equation of state parameter (\(\omega _{\phi }(z)\)) for dark energy. We also perform the data analysis for three different functional forms of \(\omega _{\phi }(z)\) by using the combination of SN Ia, BAO, and CMB datasets. We have found that for all the choices of \(\omega _{\phi }(z)\), the SN Ia \(+\) CMB/BAO dataset favors the past decelerated and recent accelerated expansion phase of the universe. Furthermore, using the combined dataset, we have observed that the reconstructed results of \(\omega _{\phi }(z)\) and q(z) are almost choice independent and the resulting cosmological scenarios are in good agreement with the \(\Lambda \)CDM model (within the \(1\sigma \) confidence contour). We have also derived the form of the potentials for each model and the resulting potentials are found to be a quartic potential for constant \(\omega _{\phi }\) and a polynomial in \(\phi \) for variable \(\omega _{\phi }\).

Keywords

Dark Energy Scalar Field Lagrangian Density Deceleration Parameter Dark Energy Model 
These keywords were added by machine and not by the authors. This process is experimental and the keywords may be updated as the learning algorithm improves.

1 Introduction

One of the biggest challenges in modern cosmology is understanding the nature of the dark energy (DE), which seems to be responsible for the observed accelerated expansion phase of the universe at the present epoch [1, 2]. Among the many candidates for DE, the cosmological constant (\(\Lambda \)) emerges as the most natural and the simplest possibility. However, \(\Lambda \)-cosmology suffers from the so-called “fine tuning” and “cosmic coincidence” problems [3, 4]. These theoretical problems motivated cosmologists to think beyond the cosmological constant and explore other unknown components which may be responsible for the late-time accelerated expansion phase of the universe. The scalar field models have played an important leading role as a candidate of DE due to its dynamical nature and simplicity. Till now, a variety of scalar field DE models have been proposed, such as quintessence (canonical scalar field), k-essence, phantom, tachyon, dilatonic dark energy, and so on (for details, see Ref. [5] and the references therein). But the origin and nature of DE still remains completely unknown, despite many years of research.

It is strongly believed that the universe had a rapid exponential expansion phase during a short era in the very early epoch. This is known as inflation [6, 7]; it can give a satisfactory explanation to the problems of the Hot Big Bang cosmology (for example, the horizon, flatness, and monopole problems). Generally, cosmologists realized this inflationary scenario by using a single canonical scalar field called the “inflaton”, which has a canonical kinetic energy term (\(\frac{{\dot{\phi }}^{2}}{2}\)) in the Lagrangian density. In the literature, there also exist some inflationary models in which the kinetic energy term is different from the standard canonical scalar field case (instead of the standard form \(\frac{{\dot{\phi }}^{2}}{2}\)). Such models are commonly known as the non-canonical scalar field models of inflation. Such non-canonical scalar fields have been found to have many attractive features compared to the canonical scalar field case, for example, the slow-roll conditions can be achieved more easily as compared to the canonical case. Many interesting possibilities with these models have been recently studied in the literature (see Refs. [8, 9, 10, 11, 12, 13, 14, 15, 16, 17, 18, 19, 20, 21, 22, 23]). It has been first shown in Refs. [8, 9] that the k-essence model (which belongs to an important class of non-canonical scalar field models) is capable of generating inflation in the early epoch. Later, Chiba et al. [10] showed that such models can equally effectively describe a DE scenario. Since the nature of DE is completely unknown, it is quite reasonable to consider a non-canonical scalar field as a candidate for DE component and check for the viability of such models. Within the framework of a non-canonical scalar field, in this work, we shall try to obtain an observationally viable cosmological model to analyze the behavior of the deceleration parameter (q) and the equation of state (EoS) parameter (\(\omega _{\phi }\)) for describing the expansion history of the universe. The motivation for this work is discussed in detail in Sect. 2. As already mentioned, as the nature of DE is unknown to us, we eventually have no firm idea regarding whether the EoS parameter of DE is a constant quantity or whether it is dynamical in nature. In this connection, the most effective choice is to assume a specific functional form for the dark energy EoS parameter \(\omega _{\phi }\) as a function of the redshift z (for details see Sect. 3.2). To study the non-canonical scalar field DE model in a more general framework, in this paper, we have considered both possibilities. First, we have studied the model for a constant EoS parameter \(\omega _{\phi }\), which is in the range \(-1<\omega _{\phi }<-\frac{1}{3}\) so as to obtain acceleration. Second, we have considered three different choices for \(\omega _{\phi }(z)\) in order to cover a wide range of the DE evolution. We have then solved the field equations and analyzed the respective cosmological scenarios for all the cases. For all the models, the deceleration parameter q is found to exhibit an evolution from early deceleration to late-time acceleration phase of the universe. This feature is essential for the structure formation of the universe. For all the models, we have also derived the potential \(V(\phi )\) in terms of the scalar field \(\phi \) by considering a specific parametrization of \(f(\phi )\). In order to compare the theoretical models of DE (for dynamical \(\omega _{\phi }\)) with the observations, we have used the SN Ia, BAO, and CMB dataset to constrain the various model parameters (for details see Appendix A). We have found that the combined dataset favors the \(\Lambda \)CDM model within the \(1\sigma \) confidence contour. We give the detailed results of this work in Sect. 4.

The present paper is organized in the following way. In the next section, we introduce some basic equations of a general non-canonical scalar field model and also discuss the motivation of this work. We then obtain the general solutions of the field equations for a particular choice of the function \(f(\phi )\) and for different forms of the EoS parameter \(\omega _{\phi }\). In Sect. 4, we summarize the results of this work. Finally, the conclusions of this work are presented in Sect. 5. Additionally, for completeness, we perform the combined data analysis in Appendix A and find the observational constraints on \(\omega _{\phi }(z)\) and q(z) using the SN Ia, BAO, and CMB datasets.

2 Basic framework

Usually, the scalar field models are characterized by a general action which has the following functional form:
$$\begin{aligned} S={\int \mathrm{d}^{4}x \sqrt{-g}{\left( \frac{R}{2} + \mathcal{L}(\phi ,X)\right) }} + S_{m} \end{aligned}$$
(1)
where R is the Ricci scalar, and \(\mathcal{L}(\phi ,X)\) is the Lagrangian density, which is an arbitrary function of the scalar field \(\phi \) and its kinetic term X. The kinetic term X is defined as \(X=\frac{1}{2}\partial _\mu \phi \partial ^\mu \phi \), which is a function of time only. The last term, \(S_{m}\), represents the action of the background matter. Throughout this paper we shall work in natural units, such that \(8\pi G=c=1\).
The expressions for the energy density (\(\rho _{\phi }\)) and pressure (\(p_{\phi }\)) associated with the scalar field are given by
$$\begin{aligned} \rho _{\phi }= & {} 2X\frac{\partial \mathcal{L}}{\partial X}- \mathcal{L}\end{aligned}$$
(2)
$$\begin{aligned} p_{\phi }= & {} \mathcal{L}(\phi ,X) \end{aligned}$$
(3)
where \(X=\frac{1}{2}{\dot{\phi }}^{2}\).
In general, the Lagrangian for a scalar field model can be represented as (Melchiorri et al. [24])
$$\begin{aligned} \mathcal{L}(\phi , X)=f(\phi )F(X)-V(\phi ) \end{aligned}$$
(4)
where \(f(\phi )\) and F(X) are arbitrary functions of \(\phi \) and X, respectively. \(V(\phi )\) is the potential for the scalar field \(\phi \).
Let us consider a homogeneous, isotropic, and spatially flat FRW universe which is characterized by the following line element:
$$\begin{aligned} \mathrm{d}s^{2}=\mathrm{d}t^{2}-a^{2}(t){\left[ \mathrm{d}r^{2} + r^{2}\mathrm{d}{\theta }^{2} + r^{2}\mathrm{sin}^{2}\theta \mathrm{d}{\phi }^{2}\right] } \end{aligned}$$
(5)
where a(t) is the scale factor of the universe. With the FRW geometry, the equations of motion take the form
$$\begin{aligned} 3H^{2}= 2f(\phi )X F_{X} -f(\phi )F + V(\phi ) + \rho _{m}, \end{aligned}$$
(6)
$$\begin{aligned} \dot{H}= - \frac{1}{2} {\left[ 2f(\phi )XF_{X} + \rho _{m} \right] }, \end{aligned}$$
(7)
$$\begin{aligned}&{\left[ f(\phi )F_{X} + 2f(\phi )XF_{XX} \right] }{\ddot{\phi }} +3Hf(\phi ){\dot{\phi }}F_{X} + 2X\frac{\partial f(\phi )}{\partial \phi }F_{X}\nonumber \\&-\frac{\partial f(\phi )}{\partial \phi }F + \frac{\partial V(\phi )}{\partial \phi } = 0, \end{aligned}$$
(8)
$$\begin{aligned} {\dot{\rho }}_{m} + 3H\rho _{m}=0, \end{aligned}$$
(9)
where \(H=\frac{\dot{a}}{a}\) denotes the Hubble parameter, an overdot indicates differentiation with respect to the time coordinate t, and \(\rho _{m}\) represents the energy density of the matter component of the universe, \(F_{X} \equiv \frac{\partial F}{\partial X}\) and \(F_{XX}\equiv \frac{\partial ^{2}F}{\partial X^{2}}\).
It deserves mentioning that Eq. (4) includes all the popular single scalar field models. It reduces to a canonical scalar field model when \(f(\phi )=\mathrm{constant}=1\) and \(F(X)=X\). Again, it describes a pure k-essence model when \(V(\phi )=0\) and a phantom scalar field model when \(f(\phi )=1\) and \(F(X)=-X\). It is interesting to note that Eq. (4) reduces to the general non-canonical scalar field model [\(\mathcal{L}(\phi , X)=F(X)-V(\phi )\)] when \(f(\phi )=1\). This type of non-canonical scalar field models was proposed by Fang et al. [21]. They studied several aspects of this type of scalar fields for different forms of F(X). Recently, these types of non-canonical scalar field models have gathered attention due to their simplicity. Unnikrishnan et al. [11] have showed that for non-canonical scalar field models, the slow-roll conditions can be more easily satisfied compared to the canonical inflationary theory. They have shown that such models (with quadratic and quartic potentials) are more consistent with the current observational constraints relative to the canonical inflation. They have also shown that such non-canonical models can drop the tensor-to-scalar ratio rather than their canonical counterparts. In fact, a lot of work has been done in the framework of the non-canonical inflationary scenario in the early epoch [8, 9, 10, 11, 12, 13, 14, 15, 16, 17, 18, 19, 20, 21]. Furthermore, Franche et al. [25] showed that the non-canonical scalar fields are the most universal case with a general Lagrangian density satisfying certain conditions. These interesting properties of non-canonical scalar field models motivated us to study the cosmological aspects of such fields in a more general framework in the context of dark energy. In the literature, a large number of functional forms of \(\mathcal{L}(\phi , X)\) have been proposed so far; see for example [11, 26, 27, 28]. In our earlier work [22, 23], we have considered a Lagrangian density of the following form:
$$\begin{aligned} \mathcal{L}(\phi , X)=X^{2} - V(\phi ), \end{aligned}$$
(10)
which can be obtained from the general form of the Lagrangian density [11, 27, 28]
$$\begin{aligned} \mathcal{L}(\phi , X)=X{\left( \frac{X}{M^{4}_{Pl}}\right) }^{\alpha -1} - V(\phi ) \end{aligned}$$
(11)
for \(\alpha =2\) and \(M_{Pl}=\frac{1}{\sqrt{8\pi G}}=1\). The above equation describes a purely canonical scalar field Lagrangian density [\(\mathcal{L}(\phi , X)=X - V(\phi )\)] when \(\alpha = 1\).
In this present work, we try to extend our previous work in [22, 23] by considering a general non-canonical scalar field model which has the following Lagrangian density:
$$\begin{aligned} \mathcal{L}(\phi , X)=f(\phi )X{\left( \frac{X}{M^{4}_{Pl}}\right) }^{\alpha -1} - V(\phi ). \end{aligned}$$
(12)
Here, following [22, 23] also, we consider \(\alpha =2\) and \(M_{Pl}=\frac{1}{\sqrt{8\pi G}}=1\).
In this case, the energy density and pressure of the scalar field are given by
$$\begin{aligned} \rho _{\phi }=\frac{3}{4}f(\phi ){\dot{\phi }}^{4} + V(\phi ) ,\end{aligned}$$
(13)
$$\begin{aligned} p_{\phi }=\frac{1}{4}f(\phi ){\dot{\phi }}^{4} - V(\phi ). \end{aligned}$$
(14)
It is evident from Eqs. (13) and (14) that the usual definition of \(\rho _{\phi }\) \((=\frac{1}{2}{\dot{\phi }}^2+V(\phi ))\) and \(p_{\phi }\) \((=\frac{1}{2}{\dot{\phi }}^2-V(\phi ))\) for a standard canonical scalar field model gets modified due to the Lagrangian density (12). Also, the equations of motion for this Lagrangian (Eqs. (6)-(9)) turn out to be
$$\begin{aligned} 3H^{2}= \frac{3}{4}f(\phi ){\dot{\phi }}^{4} + V(\phi ) + \rho _{m}, \end{aligned}$$
(15)
$$\begin{aligned} \dot{H}= - \frac{1}{2} {\left[ f(\phi ){\dot{\phi }}^{4} + \rho _{m} \right] }, \end{aligned}$$
(16)
$$\begin{aligned} {\dot{\rho }}_{\phi } + 3H(\rho _{\phi } + p_{\phi })=0, \end{aligned}$$
(17)
$$\begin{aligned} {\dot{\rho }}_{m} + 3H\rho _{m}=0. \end{aligned}$$
(18)
The solution for \(\rho _{m}\) from Eq. (18) is obtained:
$$\begin{aligned} \rho _{m}(z)=\rho _{m0}(1+z)^{3} \end{aligned}$$
(19)
where \(\rho _{m0}\) is the matter density at the present time, \(z=\frac{a_{0}}{a}-1\) is the redshift, and the present value of the scale factor \(a_{0}\) is normalized to unity.
One of the important quantities in cosmology is the dark energy EoS parameter \(\omega _{\phi }=\frac{p_{\phi }}{\rho _{\phi }}\), which, in our case, is given by
$$\begin{aligned} \omega _{\phi }=\frac{f(\phi ){\dot{\phi }}^{4} - 4V(\phi )}{3f(\phi ){\dot{\phi }}^{4} + 4V(\phi )} .\end{aligned}$$
(20)
From Eq. (17), one can then obtain the expression for the energy density of the scalar field as
$$\begin{aligned} \rho _{\phi }(z) =\rho _{\phi 0} \mathrm{exp}{\left[ 3\int ^z_{0}{\frac{1+\omega _{\phi }(z^{\prime })}{1+z^{\prime }}\mathrm{d}z^{\prime }}\right] } \end{aligned}$$
(21)
where \(\rho _{\phi 0}\) is an integration constant.
Now, the Friedmann equation can be written in the following integrated form:
$$\begin{aligned} H^2(z)=H^2_{0}{\left[ \Omega _{m0}(1+z)^3 + \Omega _{\phi 0}\mathrm{exp}{\left( 3\int ^z_{0}{\frac{1+\omega _{\phi }(z^{\prime })}{1+z^{\prime }}\mathrm{d}z^{\prime }}\right) }\right] } \end{aligned}$$
(22)
where \(H_{0}\) is the Hubble parameter at the present epoch, \(\Omega _{m0}=\frac{\rho _{m0}}{3H^{2}_{0}}\), and \(\Omega _{\phi 0}(=\frac{\rho _{\phi 0}}{3H^{2}_{0}})=1-\Omega _{m0}\) are the density parameters of the matter and scalar field (or dark energy), respectively, at the present epoch.
For this model, from Eqs. (13) and (14), the potential can be expressed (as a function of redshift z) as
$$\begin{aligned} V(z)=\frac{1}{4} (1-3\omega _{\phi }(z)) \rho _{\phi }(z) \end{aligned}$$
(23)
and
$$\begin{aligned} f(\phi ){\dot{\phi }}^4 = (1+\omega _{\phi })\rho _{\phi }. \end{aligned}$$
(24)
In order to solve the field equations analytically we will proceed as follows. Out of Eqs. (15)–(18), only three are independent equations in view of the Bianchi identity (with five unknown quantities: H, \(\rho _{m}\), \(f(\phi )\), \(V(\phi )\), and \(\phi \)). So naturally one has to assume two relationships among the different variables to solve the system of equations.
Following the above argument, in this paper, we assume that the quantity f has the functional form
$$\begin{aligned} f={\left( {\frac{f_{0}}{H}}\right) }^4 \end{aligned}$$
(25)
where \(f_{0}\) is an arbitrary constant. It deserves mentioning here that the above parametrization of \(f(\phi )\) helps us to close the system of equations. With this input, Eq. (24) can be written in the following integral form:
$$\begin{aligned} \phi (z) = \phi _{0} + \int ^z_{0} {F(z^{\prime })\mathrm{d}z^{\prime }} \end{aligned}$$
(26)
where \(\phi _{0}\) is an arbitrary constant of integration and \(F(z^{\prime }) =\frac{[(1+\omega _{\phi }(z^{\prime }))\rho _{\phi }(z^{\prime })]^{\frac{1}{4}}}{f_{0}(1+z^{\prime })}\).
Another important observable quantity, the deceleration parameter q(z), can also be expressed in terms of H(z) as
$$\begin{aligned} q(z)= -\frac{\ddot{a}}{aH^2}=-1 + (1+z) \frac{\mathrm{d}\ln {H(z)}}{\mathrm{d}z}, \end{aligned}$$
(27)
which describes the evolution of our universe.

We shall now concentrate on the dark energy EoS parameter \(\omega _{\phi }(z)\). If a function of \(\omega _{\phi }(z)\) is given, then we can find the evolution of \(\rho _{\phi }(z)\) from Eq. (21). As a result, we can also find the evolutions of H(z), q(z), V(z), and \(\phi (z)\). Inverting \(\phi (z)\) into \(z(\phi )\) and using Eq. (23), one can then obtain the potential \(V(\phi )\) in terms of \(\phi \). As already mentioned, we have considered a specific parametrization of \(f(\phi )\) and still we need another assumption to match the number of unknown parameters with the number of independent equations. With this freedom, we choose different functional forms for \(\omega _{\phi }(z)\), the equation of state parameter. In the next section, we try to obtain the functional forms of various cosmological parameters for different choices of \(\omega _{\phi }(z)\) and study their cosmological implications.

3 Theoretical models

In this section, we shall consider two phenomenological DE models for obtaining the current acceleration of the universe in the framework of a general non-canonical scalar field theory.

3.1 Model I: Accelerating universe driven by a constant EoS parameter for dark energy \((-1< \omega _{\phi } < -\frac{1}{3})\)

In this model, we shall investigate the properties of an accelerated expanding universe driven by a non-canonical scalar field dark energy with a constant EoS parameter. Recent observations suggest that the dark energy EoS \(\omega _{\phi }\) is very close to \(-1\) and the approximate bound on \(\omega _{\phi }\) is \(-1.1\le \omega _{\phi }\le -0.9\) [29, 30]. Keeping this limit in mind, we choose a constant \(\omega _{\phi }\) in the limit \(-1< \omega _{\phi } < -\frac{1}{3}\), which ensures that the model does not deviate much from a \(\Lambda \)CDM model.

In this case, the energy density of DE can be obtained (from Eq. (21)):
$$\begin{aligned} \rho _{\phi }(z)=\rho _{\phi 0}(1+z)^{3(1+\omega _{\phi })} \end{aligned}$$
(28)
where \(\rho _{\phi 0}\) is the integration constant which represents the dark energy density at the present time. Then the corresponding Friedmann equation becomes
$$\begin{aligned} H^2(z)=H^{2}_{0}{\left[ \Omega _{m0}(1+z)^{3} + (1-\Omega _{m 0})(1+z)^{3(1+\omega _{\phi })}\right] } .\end{aligned}$$
(29)
Inserting the Hubble parameter H(z) (as given in the above equation) into Eq. (27), we obtain the deceleration parameter as
$$\begin{aligned} q(z)=-1 + \frac{3 + 3\kappa _{1}(1+z)^{3\omega _{\phi }}}{2 + 2\kappa _{2}(1+z)^{3\omega _{\phi }}} \end{aligned}$$
(30)
where \(\kappa _{1}=\frac{(1-\Omega _{m 0})(1+\omega _{\phi })}{\Omega _{m 0}}\) and \(\kappa _{2}=\frac{1-\Omega _{m 0}}{\Omega _{m 0}}\). The plot of q(z) against z is shown in Fig. 1 for different values of \(\omega _{\phi }\) (within the range \(-1<\omega _{\phi }<-\frac{1}{3}\)) and \(\Omega _{m 0}=0.3\). Figure 1 shows that q(z) crosses its transition point from its positive value regime to the negative value regime in the recent past, which is consistent with the independent measurements reported by several authors (see Ref. [31] and the references therein).
Fig. 1

Plot of q vs. z for \(\omega _{\phi }=-0.7\) (thick curve), \(\omega _{\phi }=-0.8\) (dashed curve), and \(\omega _{\phi }=-0.9\) (dotted curve). For all these plots, we have taken \(\Omega _{m0}=0.3\). The horizontal line is for \(q(z)=0\)

With the help of Eqs. (26) and (28), one can solve for the scalar field in the flat FRW universe as
$$\begin{aligned} \phi (z)=\phi _{0} + \beta (1+z)^{\frac{3}{4}(1+\omega _{\phi })} \end{aligned}$$
(31)
where \(\beta =\frac{4}{3f_{0}}{\left[ \frac{3H^{2}_{0}(1-\Omega _{m0})}{(1+\omega _{\phi })^{3}}\right] }^{\frac{1}{4}}\).
With the help of Eqs. (23) and (31), we find the form of the potential in terms of \(\phi \) as
$$\begin{aligned} V(\phi )=V_{0}(\phi -\phi _{0})^4 \end{aligned}$$
(32)
where \(V_{0}=\frac{3H^2_{0}(1-\Omega _{m 0})(1-3\omega _{\phi })}{4\beta ^4}\). Thus, the constant \(\omega _{\phi }\) model leads to a quartic potential. For \(\phi _{0}=0\), the above potential is similar to the potential used by Linde [7] in the context of chaotic inflation. Using the expression for \(\phi (z)\), we also obtain the functional form of \(f(\phi )\) as
$$\begin{aligned} f(\phi )= {f_{\alpha }{\left( \phi - \phi _{0}\right) }^{-\frac{8}{1 + \omega _{\phi }}}} {{\left( 1 + f_{\beta }{\left( \phi - \phi _{0}\right) }^{\frac{4\omega _{\phi }}{1 + \omega _{\phi }}}\right) }^{-2}} \end{aligned}$$
(33)
where \(f_{\alpha }=\frac{f^{4}_{0}{\beta }^{\frac{8}{1 + \omega _{\phi }}}}{H^{4}_{0}\Omega ^{2}_{m0}}\) and \(f_{\beta }=\frac{(1-\Omega _{m0})}{\Omega _{m0}}\beta ^{-\frac{4\omega _{\phi }}{1 + \omega _{\phi }}}\). It is evident from the above equation that the function \(f(\phi )\) can be arranged as a series expansion in powers of (\(\phi - \phi _{0}\)).

3.2 Model II: Accelerating universe driven by time-dependent EoS parameter for dark energy

We shall now focus on the second possibility i.e., the EoS parameter \(\omega _{\phi }\) is dynamical in nature. For this purpose, in this subsection, we have considered three different choices of \(\omega _{\phi }\) to study the model in a more general way. If \(\omega _{\phi }\) is dynamical in nature, then one way to study models going beyond the cosmological constant is by using a particular functional form for the dark energy EoS parameter \(\omega _{\phi }(z)\). However, a large number of functional forms for \(\omega _{\phi }(z)\) have appeared in the literature [32, 33, 34, 35, 36, 37, 38, 39, 40, 41]. Usually, the parametrized form of \(\omega _{\phi }(z)\) is written as [5]
$$\begin{aligned} \omega _{\phi }(z)=\sum _{n=0} \omega _{n}x_{n}(z) \end{aligned}$$
(34)
where the \(\omega _{n}\)’s are arbitrary constants and the \(x_{n}(z)\)’s are functions of the redshift z. The numerical values of the \(\omega _{n}\)’s can be found by fitting to the observational data. Following the first order expansions in Eq. (34), several authors considered many functional forms for \(x_{n}(z)\) to investigate the evolution of \(\omega _{\phi }(z)\).
For example:
  1. (i)

    \(\omega _{\phi }(z)=\omega _{0}~=~\)constant (as we have discussed in model I) for \(x_{0}(z)=1\) and \(x_{n}=0\) (\(n\ge 1\)).

     
  2. (ii)

    \(\omega _{\phi }(z) = \omega _{0} + \omega _{1} z\) i.e., linear redshift parametrization [32, 33], for \(x_{n}(z)=z^{n}\) with \(n\le 1\).

     
  3. (iii)

    \(\omega _{\phi }(z) = \omega _{0} + \omega _{1} \mathrm{log}(1+z)\) i.e., logarithmic parametrization [36], for \(x_{n}(z)=[\mathrm{log}(1+z)]^{n}\) with \(n\le 1\).

     
  4. (iv)

    \(\omega _{\phi }(z) = \omega _{0} + \omega _{1} \frac{z}{(1+z)}\) i.e., CPL parametrization [34, 35], for \(x_{n}={\left( \frac{z}{1+z}\right) }^{n}\) with \(n=1\).

     
There are many more. It is worth mentioning that the parametrizations (ii) and (iii) diverge at high redshifts, whereas the fourth one blows up in the future, when \(z\rightarrow -1\). It should be noted that the assumed parametrization would lead to possible biases in the study of the evolution of the DE but in the absence of any information regarding the true nature of DE, these parametrizations provide some insight regarding the possible nature of DE component and are worth studying. In this paper, we shall consider two different divergence-free functional forms of \(\omega _{\phi }(z)\) which do not diverge in the future (\(z\rightarrow -1\)). In addition, we shall also consider the linear redshift parametrization of \(\omega _{\phi }(z)\) for the statistical model comparisons with the divergence-free parametrizations, at low redshifts. In order to explore the evolution of DE, we shall also try to reconstruct the deceleration parameter q(z) using Eq. (27) for these different choices in Sect. 4.

\(\bullet \) Assumption I

Here, we consider the linear redshift parametrization of the EoS parameter \(\omega _{\phi }\) [32, 33], which has the following functional form:
$$\begin{aligned} \omega _{\phi }(z) = \omega _{0} + \omega _{1} z \end{aligned}$$
(35)
where \(\omega _{0}\) represents the present value of \(\omega _{\phi }\) and the second term measures the variation of \(\omega _{\phi }\) with respect to z.
Inserting for \(\omega _{\phi }(z)\) from Eq. (35) into Eq. (21), we obtain
$$\begin{aligned} \rho _{\phi }(z)=\rho _{\phi 0}(1+z)^{(1+\omega _{0} - \omega _{1})}\mathrm{exp}(3\omega _{1}z) .\end{aligned}$$
(36)
Now Eq. (22) can be written as
$$\begin{aligned} H^2(z)=&H^2_{0}\left[ \Omega _{m0}(1+z)^3 \right. \nonumber \\&\left. + (1-\Omega _{m0})(1+z)^{(1+\omega _{0} - \omega _{1})}\mathrm{exp}(3\omega _{1}z)\right] . \end{aligned}$$
(37)
Now by numerical investigations, we have plotted V as a function of \(\phi \) in Fig. 2 by considering \(\omega _{0}=-0.95\), \(\omega _{1}=0.15\), \(\Omega _{m0}=0.3\), \(f_{0}=1\), and \(\phi _{0}=150\) for this case. Figure 2 shows that the potential \(V(\phi )\) increases initially but becomes almost flat as \(\phi \) increases. The reason behind this seems to be the form of the linear parametrization, which is appropriate only for low redshifts (\(z<<1\)) and diverges for large redshifts. The corresponding expressions for \(V(\phi )\) and \(f(\phi )\) become approximately equal to (for details see Appendix B)
$$\begin{aligned} V(\phi )\simeq 253.4{\phi }^{3} -75.04 {\phi }^{2} -512.4\phi + 10770 \end{aligned}$$
(38)
and
$$\begin{aligned} f(\phi )\simeq 2.53\times 10^{-8} \mathrm{exp}(32.17\phi ). \end{aligned}$$
(39)
Fig. 2

Plot of \(V(\phi )\) versus \(\phi \) for the linear parametrization \( \omega _{\phi }=\omega _{0} + \omega _{1}z\), by assuming \(\omega _{0}=-0.95\), \(\omega _{1}=0.15\), \(\Omega _{m0}=0.3\), \(f_{0}=1\), \(\phi _{0}=150\), and \(H_{0}=72\) km \(\mathrm{s}^{-1}~\mathrm{Mpc}^{-1}\)

In the context of DE (as it is a late-time phenomenon), the above choice of \(\omega _{\phi }(z)\) has been widely used due to its simplicity and we find for the present parametrization of f given in Eq. (25) that the potential turns out to be a polynomial in \(\phi \).

\(\bullet \) Assumption II

Next, we propose
$$\begin{aligned} \omega _{\phi }(z) = \omega _{2} + \frac{1}{1+ \frac{\omega _{3}}{(1+z)^3}} \end{aligned}$$
(40)
where \(\omega _{2}\) and \(\omega _{3}\) are arbitrary constants to be fixed by observations. It is easy to see that the EoS parameter reduces to
$$\begin{aligned} {\omega _{\phi }(z)} = \left\{ \begin{array}{ll} 1+\omega _{2},&{}{ for}\ z\rightarrow +\infty \quad (\mathrm{early~epoch}),\\ \omega _{2} + \frac{1}{1+ \omega _{3}}, &{} \text {for}~ z=0 \quad (\mathrm{present~epoch}), \\ \omega _{2},&{} \text {for} ~ z\rightarrow -1 \quad (\mathrm{far~future}). \end{array}\right. \end{aligned}$$
(41)
Thus the above choice of \(\omega _{\phi }(z)\) is a bounded function of the redshift throughout the entire cosmic evolution and it also overcomes the shortcomings of the linear and CPL parametrizations of \(\omega _{\phi }(z)\). Although this is the main motivation of proposing the ansatz given in Eq. (40), it can also be thought of as a particular form of Eq. (34) for appropriate choices of the \(\omega _{n}\)’s and \(x_{n}(z)\)’s.
Fig. 3

Plot of \(V(\phi )\) vs. \(\phi \) for the ansatz given by Eq. (40). The plot is for \(\omega _{2}=-1.25\), \(\omega _{3}=5\), \(\Omega _{m0}=0.3\), \(f_{0}=1\), \(\phi _{0}=150\), and \(H_{0}=72\) km \(\mathrm{s}^{-1}~\mathrm{Mpc}^{-1}\)

Fig. 4

The variation of \(V(\phi )\) with \(\phi \) for the ansatz given by Eq. (46). The plot is for \(A_{0}=3.5\), \(A_{1}=0.2\), \(A_{2}=0.4\), \(\Omega _{m0}=0.3\), \(f_{0}=1\), \(\phi _{0}=150\), and \(H_{0}=72\) km \(\mathrm{s}^{-1} ~\mathrm{Mpc}^{-1}\)

Fig. 5

This figure shows the \(1\sigma \) and \(2\sigma \) confidence contours for each choice of \(\omega _{\phi }(z)\) using the SN Ia \(+\) BAO/CMB dataset. The plots are for \(\Omega _{m0}=0.3\) (for choices I and II) and \(\Omega _{m0}=0.3\), \(A_{0}=3.5\) for choice III. The upper, middle and lower panels represent the \(\omega _{0}-\omega _{1}\), \(\omega _{2}-\omega _{3}\), and \(A_{1}-A_{2}\) parameter space for the choices I, II, and III, respectively. In each panel, the large dot represents the best-fit values of the model parameters, whereas the small dot represents the chosen values of these parameters in the analytical models (as mentioned in Sect. 3.2). The corresponding \(\chi ^{2}\) for the best-fit points are displayed in the Table 1

In this case, \(\rho _{\phi }(z)\) and H(z) evolve as
$$\begin{aligned} \rho _{\phi }(z)=\frac{\rho _{\phi 0}}{1+\omega _{3}}(1+z)^{3(1+\omega _{2})}(\omega _{3} + (1+z)^3) \end{aligned}$$
(42)
where \(\rho _{\phi 0}\) is the present value of the scalar field density, and
$$\begin{aligned} H^2(z)&=H^{2}_{0}[\Omega _{m0}(1+z)^{3}\nonumber \\&\quad + \frac{(1-\Omega _{m0})}{1+\omega _{3}}(1+z)^{3(1+\omega _{2})}(\omega _{3} + (1+z)^3)]. \end{aligned}$$
(43)
For this specific choice, \(V(\phi )\) and \(f(\phi )\) can be obtained (see Appendix B):
$$\begin{aligned} V(\phi )\simeq 0.11\phi ^{4} -58.5\phi ^{3} + 12136\phi ^{2} -10^{6}\phi + 4\times 10^{7}\nonumber \\ \end{aligned}$$
(44)
and
$$\begin{aligned} f(\phi )\simeq f_{1}\phi ^{6} + f_{2}\phi ^{5}+f_{3}\phi ^{4} + f_{4}\phi ^{3} + f_{5}\phi ^{2} + f_{6}\phi + f_{7}\nonumber \\ \end{aligned}$$
(45)
where \(f_{1}=6\times 10^{-15}\), \(f_{2}=-5\times 10^{-12}\), \(f_{3}=10^{-9}\), \(f_{4}=-2\times 10^{-7}\), \(f_{5}=2\times 10^{-5}\), \(f_{6}=-0.001\), and \(f_{7}=0.0205\). These values of \(f_{n}\)’s have been obtained for \(\omega _{2}=-1.25\), \(\omega _{3}=5\), \(\Omega _{m0}=0.3\), \(f_{0}=1\), and \(\phi _{0}=150\). In this case, the evolution of the potential \(V(\phi )\) is shown in Fig. 3 and we have seen that \(V(\phi )\) sharply decreases with \(\phi \) from an extremely large value to a fixed value.
Table 1

Best-fit values for various model parameters for the analysis of SN Ia \(+\) BAO/CMB dataset. Here, \(\omega _{\phi }(z=0)\) represents the present value of the EoS parameter \(\omega _{\phi }(z)\) in the best-fit models. For this analysis, we have considered \(\Omega _{m0}=0.3\) (for choices I and II) and \(\Omega _{m0}=0.3\), \(A_{0}=3.5\) for choice III

Choice

Best-fit values of

\(\omega _{\phi }(z=0)\)

\({\chi ^{2}_{m}}\)

 

model parameters

  

I

\(\omega _{0}=-1.01\)

\(\omega _{1}=0.10\)

\(-1.01\)

599.90

II

\(\omega _{2}=-1.19\)

\(\omega _{3}=7\)

\(-1.06\)

565.43

III

\(A_{1}=-0.16\)

\(A_{2}=-0.14\)

\(-1.04\)

564.86

\(\bullet \) Assumption III

The next choice adopted in this paper was suggested by Alam et al. [39, 40], which is the functional form
$$\begin{aligned} \omega _{\phi }(z)=-1 + \frac{A_{1}(1+z) + 2A_{2}(1+z)^2}{3[A_{0} +A_{1}(1+z) + A_{2}(1+z)^2]}. \end{aligned}$$
(46)
This choice is exact and gives the cosmological constant \(\omega _{\phi }=-1\) for \(A_{1}=A_{2}=0\) and DE models with \(\omega _{\phi }=-\frac{1}{3}\) for \(A_{0}=A_{1}=0\) and \(\omega _{\phi }=-\frac{2}{3}\) for \(A_{0}=A_{2}=0\). The above choice mimics a DE model very well, and also it can be viewed as a power law in the redshift dependence of the energy density for DE component. With this choice of \(\omega _{\phi }(z)\), Eq. (21) immediately gives
$$\begin{aligned} \rho _{\phi }(z)=\frac{\rho _{\phi 0}}{A_{0} +A_{1} + A_{2}} {\left[ A_{0} +A_{1}(1+z) + A_{2}(1+z)^2\right] } \end{aligned}$$
(47)
where \(\rho _{\phi 0}\) is the present value of \(\rho _{\phi }\).
Fig. 6

This figure shows the evolution of q(z) with redshift z for the choices I (upper panel), II (middle panel), and III (lower panel), respectively. The reconstruction is done using the SN Ia \(+\) BAO/CMB dataset by assuming \(\Omega _{m0}=0.3\) for choices I, II and \(\Omega _{m0}=0.3\), \(A_{0}=3.5\) for choice III. In each panel, the thick solid line shows the best-fit curve, the dashed lines represent the \(1\sigma \) confidence level, and the thin lines represent the \(2\sigma \) confidence level around the best fit. Also, in each panel, the horizontal line indicates \(q(z)=0\)

Fig. 7

The upper and middle panel represent the plot of \(\omega _{\phi }(z)\) versus z with \(\Omega _{m0}=0.3\) for ansatzes I and II, respectively. The lower panel corresponds to the evolution of \(\omega _{\phi }(z)\) for the ansatz III. This plot is for \(\Omega _{m0}=0.3\) and \(A_{0}=3.5\). Also, in each panel, the inset diagram shows the evolution of the total EoS parameter \(\omega _\mathrm{tot}(z)\) with z for these ansatzes. The thick solid line shows the best-fit curve, the dashed lines represent the \(1\sigma \) confidence level, and the thin lines represent the \(2\sigma \) confidence level around the best fit

In this case, the Hubble parameter is expressed as
$$\begin{aligned} H^2(z)&=H^{2}_{0}\left[ \phantom {\left. + \frac{(1-\Omega _{m0})}{A_{0} +A_{1} + A_{2}} {\left( A_{0} +A_{1}(1+z) + A_{2}(1+z)^2\right) }\right] }\Omega _{m0}(1+z)^{3}\right. \nonumber \\&\quad \left. + \frac{(1-\Omega _{m0})}{A_{0} +A_{1} + A_{2}} {\left( A_{0} +A_{1}(1+z) + A_{2}(1+z)^2\right) }\right] . \end{aligned}$$
(48)
We have then solved Eqs. (23) and (26) numerically and have plotted V as a function of \(\phi \) for some specific values of the model parameters (\(A_{0}=3.5\), \(A_{1}=0.2\), \(A_{2}=0.4\), \(\Omega _{m0}=0.3\), \(f_{0}=1\), and \(\phi _{0}=150\)) in Fig. 4. It is evident from Fig. 4 that the potential \(V(\phi )\) always decreases with the scalar field \(\phi \). For the present model, \(V(\phi )\) and \(f(\phi )\) can be explicitly expressed in terms of \(\phi \) as (see Appendix B)
$$\begin{aligned} V(\phi )\simeq 20920 - 13210\phi - 17300\phi ^{2} - 16510\phi ^{3} \end{aligned}$$
(49)
and
$$\begin{aligned} f(\phi )\simeq 1.004\times 10^{-53} \phi ^{21.23}. \end{aligned}$$
(50)
However, in general for Model II (which includes ansatz I, II, and III) with the particular choice of Eq. (25), one can write the potential V as a polynomial in \(\phi \) in the following manner:
$$\begin{aligned} V(\phi ) = \sum ^{n}_{i=0}V_{i}\phi ^{i} \end{aligned}$$
(51)
where \(n>0\), the \(V_{i}\)’s are constants, and the values of these parameters are different for different choices of \(\omega _{\phi }(z)\). Interestingly, we have found that it is a generalization of other well-known potentials (see [5] and the references therein), for example, a constant potential or a power-law potential. We have also found that the parametrization (25) leads to the quantity \(f(\phi )\) as an exponential, polynomial, and power law in \(\phi \) for choices I, II, and III, respectively. In the following section, we shall use these choices to discuss the possibility of constraining \(\omega _{\phi }(z)\) and q(z) from observations.

4 Results

Following the statistical analysis (see Appendix A), in this section, we present the fitting results for different choices of the EoS parameter for DE. Figure 5 shows the \(1\sigma \) and \(2\sigma \) confidence contours for each choice (I, II, and III) using the SN Ia \(+\) BAO/CMB dataset.

The best-fit values of the model parameters and \(\omega _{\phi }(z=0)\) for these different choices are given in Table 1.

Using those best-fit values, we have reconstructed the deceleration parameter q(z) for each model and the results are plotted in Fig. 6.

It is evident from Fig. 6 that q(z) shows a smooth transition from a decelerated (\(q>0\)) to an accelerated (\(q<0\)) phase of expansion of the universe at the transition redshift \(z_{t}=0.38\) (for ansatz I), 0.36 (for ansatz II), and 0.43 (for ansatz III) for the best-fit models. These results are in good agreement with those obtained by several authors based on various other considerations [42, 43, 44].

Furthermore, we also show the reconstructed evolution history of the EoS parameter in Fig. 7 for each choice of \(\omega _{\phi }(z)\).

We have also plotted the total EoS parameter, which is defined as \(\omega _\mathrm{tot}(z)=\frac{p_{\phi }}{\rho _{\phi } + \rho _{m}}\), as a function of z for these choices (see the inset diagram of Fig. 7). From Table 1, we have found that the current values of \(\omega _{\phi }(z)\) for the best-fit DE models are very close to \(-1\), i.e., the models do not deviate very far from the \(\Lambda \)CDM model (\(\omega _{\Lambda }=-1\)) at the present epoch. However, as indicated in Table 1, the present parametrized model favors a phantom model (\(\omega _{\phi }<-1\)) in the \(2\sigma \) limit and thus requires further attention.

5 Conclusions

In this work, we have studied various non-canonical scalar field DE models in a spatially flat, homogeneous, and isotropic FRW space-time. In this framework, we have obtained the general solutions of the field equations for different choices of the EoS parameter. For completeness, we have also investigated how the joint analysis of the SN Ia \(+\) BAO/CMB datasets constrains the redshift evolutions of q(z) and \(\omega _{\phi }(z)\) for different choices of \(\omega _{\phi }(z)\) (as given in Model II). In Fig. 5, we have also shown the \(1\sigma \) and \(2\sigma \) contour plots of the pairs (\(\omega _{0} , \omega _{1}\)) (upper panel), (\(\omega _{2} , \omega _{3}\)) (middle panel), and (\(A_{1} , A_{2}\)) (lower panel) for the ansatzes I, II, and III, respectively. In this analysis, we have also calculated the best-fit values of the free parameters (as shown by large dots in Fig. 5) and it has been found that the chosen values of these parameters (which were chosen for solving the parametric relations in Appendix B) are well fitted within the \(1\sigma \) confidence contour (as shown by small dots in Fig. 5).

We have shown that the deceleration parameter q undergoes a smooth transition from its deceleration phase (\(q>0\), at high z) to an acceleration phase (\(q<0\), at low z) for all of the considered parametrized models. However, as mentioned in the previous section, the value of \(z_{t}\), where the signature flip of q (from the decelerating to an accelerating expansion phase) takes place has been calculated and the results obtained are consistent with the present day cosmological observations. From the SN Ia \(+\) BAO/CMB analysis, we have also found \(q(z=0)=-0.56\), \(-0.64\), and \(-0.60\) for ansatzes I, II, and III, respectively, which also agree very well with the recent observational results.

From Table 1, we have observed that the EoS parameter \(\omega _{\phi }(z=0)\approx -1\), but slightly less than \(-1\) for all three choices (as discussed in Sect. 4). As we have seen \(\omega _{\phi }(z=0)\approx -1\), our models do not deviate very far from the \(\Lambda \)CDM model (see also Fig. 7), which is currently known as the standard model for modern cosmology. In order to gain more physical insight into these time evolutions of the EoS parameter, we have also plotted the reconstructed total EoS parameter \(\omega _\mathrm{tot}(z)\) in Fig. 7 (see the inset diagram of Fig. 7). For each choice, this figure shows that \(\omega _\mathrm{tot}(z)\) attains the required value of \(-\frac{1}{3}\) around \(z=0.62\) (within \(1\sigma \) confidence level) and remains always greater than \(-1\) up to the present epoch. These scenarios also agree very well with the observational data.

However, the models presented here are restricted because the form of \(f(\phi )\) chosen was ad hoc (as given in Eq. (25)) and did not follow from any principle. In this regard, we have mentioned earlier that we make this choice in order to close the system of equations. With this choice of \(f(\phi )\), we have derived the form of the potential \(V(\phi )\) in terms of \(\phi \) for different models. We have found that Model I leads to a quartic potential, whereas Model II leads to a polynomial potential for each choice of \(\omega _{\phi }(z)\). We have seen that, with a suitable choice of \(V_{i}\)’s for the potential (as given in Eq. (51)), it is possible to reproduce the other well-known potentials in the context of DE. However, many possibilities are opened up to accommodate a physically viable potential for other parametrizations of \(f(\phi )\) or f(H). Finally, we would like to emphasize that all the considered models provide a deceleration for high redshift and an acceleration for low redshift as required for the structure formation of the universe. However, these results are completely independent of any choice of \(f(\phi )\). With the increase of more good quality observational data at low, intermediate, and high redshifts, the constraints on \(z_{t}\) (or q(z)) and \(\omega _{\phi }(z)\) are expected to get improved in the near future.

Notes

Acknowledgments

One of the authors (AAM) is thankful to Govt. of India for financial support through Maulana Azad National Fellowship. SD wishes to thank IUCAA, Pune, for an associateship program.

References

  1. 1.
    A.G. Riess et al., Astron. J. 116, 1009 (1998)CrossRefADSGoogle Scholar
  2. 2.
    S. Perlmutter et al., Astrophys. J. 517, 565 (1999)CrossRefADSGoogle Scholar
  3. 3.
    S. Weinberg, Rev. Mod. Phys. 61, 1 (1989)MathSciNetCrossRefADSGoogle Scholar
  4. 4.
    S. Carroll, Living Rev. Relat. 4, 1 (2001)ADSGoogle Scholar
  5. 5.
    E.J. Copeland, M. Sami, S. Tsujikawa, IJMP D 15, 1753 (2006)MathSciNetCrossRefADSGoogle Scholar
  6. 6.
    A.H. Guth, Phys. Rev. D 23, 347 (1981)CrossRefADSGoogle Scholar
  7. 7.
    A.D. Linde, Phys. Lett. B 129, 177 (1983)MathSciNetCrossRefADSGoogle Scholar
  8. 8.
    C. Armendariz-Picon, T. Damour, V. Mukhanov, Phys. Lett. B 458, 209 (1999)MathSciNetCrossRefADSGoogle Scholar
  9. 9.
    J. Garriga, V.F. Mukhanov, Phys. Lett. B 458, 219 (1999)MathSciNetCrossRefADSGoogle Scholar
  10. 10.
    T. Chiba, T. Okabe, M. Yamaguchi, Phys. Rev. D 62, 023511 (2000)Google Scholar
  11. 11.
    S. Unnikrishnan et al., JCAP 018, 1208 (2012). arXiv:1205.0786 [astro-ph.CO]Google Scholar
  12. 12.
    M. Fairbairn, M.H.G. Tytgat, Phys. Lett. B 546, 1 (2002)CrossRefADSGoogle Scholar
  13. 13.
    D.A. Steer, F. Vernizzi, Phys. Rev. D 70, 043527 (2004)MathSciNetCrossRefADSGoogle Scholar
  14. 14.
    L.P. Chimento, Phys. Rev. D 69, 123517 (2004)MathSciNetCrossRefADSGoogle Scholar
  15. 15.
    R.J. Scherrer, Phys. Rev. Lett. 93, 011301 (2004)CrossRefADSGoogle Scholar
  16. 16.
    D. Bertacca, S. Matarrese, M. Pietroni, Mod. Phys. Lett. A 22, 2893 (2007)Google Scholar
  17. 17.
    G. Panotopoulos, Phys. Rev. D 76, 127302 (2007)Google Scholar
  18. 18.
    N. Bose, A.S. Majumdar, Phys. Rev. D 80, 103508 (2009)CrossRefADSGoogle Scholar
  19. 19.
    J. De-Santiago, J.L. Cervantes-Cota, Phys. Rev. D 83, 063502 (2011)CrossRefADSGoogle Scholar
  20. 20.
    T. Golanbari et al., Phys. Rev. D 89, 103529 (2014)Google Scholar
  21. 21.
    W. Fang et al., Class. Quant. Grav. 24, 3799 (2007)CrossRefADSGoogle Scholar
  22. 22.
    S. Das, A.A. Mamon, Astrophys. Space Sci. 355, 371 (2015). arXiv:1407.1666 [gr-qc]CrossRefADSGoogle Scholar
  23. 23.
    A.A. Mamon, S. Das, Eur. Phys. J. C 75, 244 (2015). arXiv:1503.06280 [gr-qc]CrossRefADSGoogle Scholar
  24. 24.
    A. Melchiorri et al., Phys. Rev. D 68, 043509 (2003)CrossRefADSGoogle Scholar
  25. 25.
    P. Franche et al., Phys. Rev. D 81, 123526 (2010)CrossRefADSGoogle Scholar
  26. 26.
    C. Armendariz-Picon, E.A. Lim, J. Cosmol. Astropart. Phys. 0508, 007 (2005)CrossRefADSGoogle Scholar
  27. 27.
    V. Mukhanov, A. Vikman, J. Cosmol. Astropart. Phys. 0602, 004 (2006)CrossRefADSGoogle Scholar
  28. 28.
    S. Unnikrishnan, Phys. Rev. D 78, 063007 (2008)CrossRefADSGoogle Scholar
  29. 29.
    W.M. Wood-Vasey et al., Astrophys. J. 666, 694 (2007)CrossRefADSGoogle Scholar
  30. 30.
    T.M. Davis et al., Astrophys. J. 666, 716 (2007)CrossRefADSGoogle Scholar
  31. 31.
    A. A. Mamon, S. Das, Int. J. Mod. Phys. D 25, 1650032 (2016). arXiv:1507.00531
  32. 32.
    D. Huterer, M.S. Turner, Phys. Rev. D 60, 081301 (1999)CrossRefADSGoogle Scholar
  33. 33.
    J. Weller, A. Albrecht, Phys. Rev. Lett. 86, 1939 (2001)CrossRefADSGoogle Scholar
  34. 34.
    M. Chevallier, D. Polarski, Int. J. Mod. Phys. D 10, 213 (2001)CrossRefADSGoogle Scholar
  35. 35.
    E.V. Linder, Phys. Rev. Lett. 90, 091301 (2003)CrossRefADSGoogle Scholar
  36. 36.
    G. Efstathiou, MNRAS 342, 810 (2000)Google Scholar
  37. 37.
    E.M. Barboza, J.S. Alcaniz, Phys. Lett. B 666, 415 (2008)CrossRefADSGoogle Scholar
  38. 38.
    C.-J. Feng et al., arXiv:1206.0063 [astro-ph.CO]
  39. 39.
    U. Alam, V. Sahni, T.D. Saini, A.A. Starobinski, MNRAS 354, 275 (2004)CrossRefADSGoogle Scholar
  40. 40.
    U. Alam, V. Sahni, A.A. Starobinski, JCAP 0406, 008 (2004)CrossRefADSGoogle Scholar
  41. 41.
    J. Weller, A. Albrecht, Phys. Rev. D 65, 103512 (2002)CrossRefADSGoogle Scholar
  42. 42.
    A.G. Riess et al., Astrophys. J. 607, 665 (2004)CrossRefADSGoogle Scholar
  43. 43.
    A.G. Riess et al., ApJ 659, 98 (2007)CrossRefADSGoogle Scholar
  44. 44.
    J.V. Cunha, J.A.S. Lima, arXiv:0805.1261 [astro-ph]
  45. 45.
    N. Suzuki et al., Astrophys. J. 746, 85 (2012). arXiv:1105.3470 [astro-ph.CO]CrossRefADSGoogle Scholar
  46. 46.
    S. Nesseris, L. Perivolaropoulos, Phys. Rev. D 72, 123519 (2005)CrossRefADSGoogle Scholar
  47. 47.
    F. Beutler et al., Mon. Not. R. Astron. Soc. 416, 3017 (2011)CrossRefADSGoogle Scholar
  48. 48.
    W.J. Percival et al., Mon. Not. R. Astron. Soc. 401, 2148 (2010)CrossRefADSGoogle Scholar
  49. 49.
    C. Blake et al., Mon. Not. R. Astron. Soc. 418, 1707 (2011)CrossRefADSGoogle Scholar
  50. 50.
    N. Jarosik et al., Astrophys. J. Suppl. 192, 14 (2011)CrossRefADSGoogle Scholar
  51. 51.
    R. Goistri et al., JCAP 03, 027 (2012)CrossRefADSGoogle Scholar

Copyright information

© The Author(s) 2016

Open AccessThis article is distributed under the terms of the Creative Commons Attribution 4.0 International License (http://creativecommons.org/licenses/by/4.0/), which permits unrestricted use, distribution, and reproduction in any medium, provided you give appropriate credit to the original author(s) and the source, provide a link to the Creative Commons license, and indicate if changes were made.

Funded by SCOAP3

Authors and Affiliations

  1. 1.Department of PhysicsVisva-BharatiSantiniketanIndia

Personalised recommendations