Models for Little Rip Dark Energy Paul H. Frampton1∗, Kevin J. Ludwick1†, Shin’ichi Nojiri2‡, Sergei D. Odintsov2,3§,¶, and Robert J. Scherrer,4∗∗ 1Department of Physics & Astronomy, University of North Carolina, Chapel Hill, NC 27599 2Department of Physics, Nagoya University, Nagoya 464-8602, Japan and Kobayashi-Maskawa Institute for the Origin of Particles and the Universe, Nagoya University, Nagoya 464-8602, Japan 3Instituciò Catalana de Recerca i Estudis Avançats (ICREA) and Institut de Ciencies de l’Espai (IEEC-CSIC), Campus UAB, Facultat de Ciencies, Torre C5-Par-2a pl, E-08193 Bellaterra (Barcelona), Spain and 4Department of Physics & Astronomy, Vanderbilt University, Nashville, TN 37235 (Dated: January 17, 2012) We examine in more detail specific models which yield a little rip cosmology, i.e., a universe in whichthedarkenergy densityincreases without boundbuttheuniverseneverreachesafinite-time singularity. Wederivetheconditionsforthelittleripintermsoftheinertialforceintheexpanding 2 universe and present two representative models to illustrate in more detail the difference between 1 littleripmodelsandthosewhichareasymptoticallydeSitter. Wederiveconditionsontheequation 0 of state parameter of the dark energy to distinguish between the two types of models. We show 2 that couplingbetween dark matterand dark energy with a little rip equation of state can alter the evolution, changing the little rip into an asymptotic de Sitter expansion. We give conditions on n minimally-coupled phantom scalar field models and on scalar-tensor models that indicate whether a or not they correspond to a little rip expansion. We show that, counterintuitively, despite local J instability, a little-rip cosmology has an infinite lifetime. 6 1 PACSnumbers: 95.36.+x,98.80.Cq ] h t I. INTRODUCTION - p e Thecurrentaccelerationoftheuniverseisoftenattributedtodarkenergy,anunknownfluidwitheffectiveequation h of state (EoS) parameter w close to 1. The observational data [1] favor ΛCDM with w = 1. However, phantom [ − − (w < 1) or quintessence ( 1/3>w > 1) dark energy models are not excluded by observational data [2]. In both 2 cases,−it is known that the−universe may−evolve to a finite-time future singularity. Phantom dark energy models can v leadtoasingularityinwhichthescalefactoranddensitybecomeinfinite atafinitetime;suchasingularityiscalleda 7 bigrip[3,4],orTypeIsingularity[5]. Forquintessencedarkenergy,onecanhaveasingularityforwhichthe pressure 6 goesto infinity ata fixedtime, butthe scale factoranddensity remainfinite; this is calledasuddensingularity[6, 7], 0 0 or a Type II singularity [5]). Alternately, the density and pressure can both become infinite with a finite scale factor . at a finite time (a Type III singularity), or higher derivatives of the Hubble parameter H can diverge (a Type IV 8 singularity)[5]. Theoccurrenceofasingularityatafinitetimeinthefuturemayleadtosomeinconsistencies. Several 0 1 scenariosto avoidafuture singularityhavebeenproposedsofar: couplingwithdarkmatter[8], inclusionofquantum 1 effects [9], additional changes in the equation of state [10] or special forms of modified gravity [10]. : Recently, a new scenario to avoid a future singularity has been proposed in Ref. [11]. In this scenario, w is less v than 1, so that the dark energy density increases with time, but w approaches 1 asymptotically and sufficiently i X rapidl−y that a singularity is avoided. This proposed non-singular cosmology was c−alled a “little rip” because it leads r to a dissolution of bound structures at some point in the future (similar to the effect of a big rip singularity). It can a be realizedintermsofageneralfluidwithacomplicatedEoS[5,12]. Theevolutionofthelittle ripcosmologyisclose to that of ΛCDM up to the present, and is similarly consistent with the observational data. The presentarticle is devoted to further study of the properties of the little rip cosmology. In the next section, the inertial force interpretation of the little rip is developed, and it becomes clear why a dissolution of bound structures occurs. Coupling of the little rip fluid with dark matter is considered in Section III. It is shown that as the result of such a coupling an asymptotically de Sitter universe can eventually evolve to have a little or big rip. In Section IV, the little rip cosmology is reconstructed in terms of scalar field models. Our results are summarized in Section V. ∗ [email protected] † [email protected] ‡ [email protected] § [email protected] ¶ AlsoatTomskStatePedagogical University. ∗∗ [email protected] 2 II. INERTIAL FORCE INTERPRETATION OF THE LITTLE RIP As the universe expands, the relative acceleration between two points separated by a comoving distance l is given by la¨/a, where a is the scale factor. An observer a comoving distance l away from a mass m will measure an inertial force on the mass of F =mla¨/a=ml H˙ +H2 . (1) iner (cid:16) (cid:17) Let us assume the two particles are bound by a constant force F . If F is positive and greater than F , the two 0 iner 0 particles become unbound. This is the “rip” produced by the accelerating expansion. Note that equation (1) shows that a rip always occurs when either H diverges or H˙ diverges (assuming H˙ > 0). The first case corresponds to a “big rip” [13], while if H is finite, but H˙ diverges with H˙ >0, we have a Type II or “sudden future” singularity [5–7], whichalso leads to a rip. However,as notedin Ref. [11], itis possible for H, andtherefore, F , to increasewithout iner bound and yet not produce a future singularity at a finite time; this is the little rip. Both the big rip and little rip are characterized by F ; the difference is that F occurs at a finite time for a big rip and as t iner iner → ∞ → ∞ → ∞ for the little rip. An interesting case occurs when H is finite and H˙ diverges but is negative. In this case, even though the universe is expanding, all structures are crushed rather than ripped. An example is given by H =H +H (t t)α . (2) 0 1 c − Here H and H are positive constants and α is a constant with 0<α<1. 0 1 By using the FRW equations 3 1 H2 =ρ, 2H˙ +3H2 =p, (3) κ2 −κ2 (cid:16) (cid:17) we may rewrite (1) in the following form: mlκ2 F = (ρ+3p) . (4) iner − 6 Here κ2 = 8πG and G is Newton’s gravitational constant. Not surprisingly, we see that the inertial force is sourced by the quantity ρ+3p. Then if we consider the general equation of state, p= ρ+f(ρ), (5) − we find mlκ2 F = (2ρ 3f(ρ)) . (6) iner 6 − AsnotedinRef.[11],whenw 1butw < 1,aripcanoccurwithoutasingularity. Ifweignorethecontribution →− − from matter, the equation of state (EoS) parameter w of the dark energy can be expressed in terms of the Hubble rate H as 2H˙ w = 1 . (7) − − 3H2 Then if H˙ >0, we find w< 1. − Now consider the following example: H =H eλt. (8) 0 Here H and λ are positive constants. Eq. (8) tells us that there is no curvature singularity for finite t. By using 0 Eq. (7), we find 2λ w = 1 e−λt, (9) − − 3H 0 and therefore w < 1 and w 1 when t + , and w is always less than 1 when H˙ is positive. From Eq. (1), − →− → ∞ − we have F =ml λH eλt+H2e2λt , (10) iner 0 0 (cid:0) (cid:1) 3 which is positive and unbounded. Thus, F becomes arbitrarily large with increasing t, resulting in a little rip. iner As another example, consider the model: H =H H e−λt. (11) 0 1 − Here H , H , and λ are positive constants and we assume H >H and t>0. Since the second term decreaseswhen 0 1 0 1 t increases, the universe goes to asymptotically de Sitter space-time. Then from Eq. (7), we find 2λH e−λt 1 w = 1 . (12) − − 3(H H e−λt)2 0 1 − As in the previous example, w < 1 and w 1 when t + . For H given by Eq. (11), however, the inertial − → − → ∞ force, given by (1), is F =ml λH e−λt+ H H e−λt 2 , (13) iner 1 0 1 − n (cid:0) (cid:1) o which is positive but bounded and F mlH2 when t + . Therefore if we choose H , H , and λ small iner → 0 → ∞ 0 1 enough, we do not obtain a rip. When t becomes large, the scale factor a is given by that of the de Sitter space-time a a eH0t, and the energy density ρ has the following form: 0 ∼ − λ ρ= 3 H2 3 H2 2H H e−λt 3 H2 2H H a H0 , (14) κ2 ∼ κ2 0 − 0 1 ∼ κ2 0 − 0 1(cid:18)a0(cid:19) ! (cid:0) (cid:1) which is an increasing function of a and becomes finite as a . →∞ For t , Eq. (12) gives the asymptotic behavior of w to be →∞ 2λH e−λt 1 w 1 , (15) ∼− − 3H2 0 which is identical with (9) if we replace λH /H with λ. 1 0 These results indicate that knowledge of the asymptotic (t ) behavior of w(t) is insufficient to distinguish → ∞ models with a rip from models which are asymptotically de Sitter. The reasonfor this becomes clear when we derive the expression for ρ(t) as a function of w(t). The evolution of ρ is given by: dρ = 3H(ρ+p), (16) dt − which can be expressed as dρ ρ−3/2 = √3κ(1+w). (17) dt − Integrating between initial and final times t and t gives: i f √3 tf ρ−1/2 ρ−1/2 = [1+w(t)]dt. (18) i − f − 2 Zti Evolution leading to a little rip implies that ρ as t , while asymptotic de Sitter evolution requires f f → ∞ → ∞ ρ constant as t . However,in either case, the integral on the right-hand side simply approaches a constant f f → →∞ as the upper limit goes to infinity. Thus, the asymptotic functionalformfor w(t) is nota goodtest ofthe asymptotic behavior of ρ. On the other hand, expressing the equation of state parameter as a function of the scale factor a instead of the time t does provide a clearer test of the existence of a future rip. Equation (16) can be written in terms of the scale factor as adρ = 3[1+w(a)], (19) ρda − from which it follows that ρ af da f ln = 3 [1+w(a)] . (20) ρ − a (cid:18) i(cid:19) Zai 4 Thus, ρ is asymptotically constant if the integral of (1+w)/a converges at its upper limit, while ρ will increase without bound, leading to a rip, when the integral diverges. Then if 1+w(a) behaves as an inverse power of a, as in 1+w(a) a−ǫ with arbitrary positive constant ǫ when a , the integration on the right-hand side of (20) is ∼ → ∞ finite when a , and therefore a rip does not occur. If 1+w(a) vanishes more slowly than any power of a when f →∞ a , e.g., 1+w(a) 1/lna, the integration on the right-hand side of (20) diverges when a , and therefore f →∞ ∼ →∞ a rip is generated. We nowconsiderwhatkindofperfectfluidrealizesthe evolutionofH inEqs.(8)or(11). The FRWequationsgive 3 2 ρ= H2, ρ+p= H˙ . (21) κ2 −κ2 Consider first the model given by Eq. (8). By substituting Eq. (8) into Eq. (21) and eliminating t, we obtain: 4λ2 (ρ+p)2 = ρ. (22) 3κ2 On the other hand, for the case corresponding to Eq. (11), we obtain: 3H2 3H 3κ2 ρ= 0 + 0 (ρ+p)+ (ρ+p)2 . (23) κ2 λ 4λ2 III. COUPLING WITH DARK MATTER In Ref. [8], it was shown that the coupling of zero-pressure dark matter with phantom dark energy could avoid a big rip singularity, and the universe might evolve to asymptotic de Sitter space. Here we investigate the possibility that coupling with the dark matter could avoid a little rip. We consider the equation of state Eq. (22), for which a little rip occurs in the absence of such a coupling. We show that by adding a coupling with dark matter, a little rip can be avoided, and the universe can evolve to de Sitter space. We now consider the following conservation law [8] ρ˙+3H(ρ+p)= Qρ, ρ˙ +3Hρ =Qρ. (24) DM DM − Here ρ is the energy density of the dark matter and Q is a positive constant. The right-hand sides in Eqs. (24) DM express the decay of the dark energy into dark matter. We assume the equation of state given in Eq. (22), for which a rip could occur. Then the first equation in (24) can be rewritten as 2λ√3ρ ρ˙ H = Qρ. (25) − κ − Note that ρ+p<0 since we are considering the model w < 1. − We now assume the de Sitter solution where H is a constant: H = H > 0. If we neglect the contribution from 0 everything other than the dark energy and dark matter, the first FRW equation 3 H2 =ρ+ρ , (26) κ2 DM indicates that ρ+ρ is a constant. Then Eq. (24) becomes DM 0=3H (ρ+p+ρ ) . (27) 0 DM Since H =H >0, we find 0 ρ = ρ p. (28) DM − − Note that the above equation (28) can be obtained from the conservation law (24) and the first FRW equation (26) without using any equation of state. Now we assume the equation of state (22). Combining Eqs. (22) and (28), we get 3κ2 ρ= ρ2 . (29) 4λ2 DM 5 Since ρ+ρ is a constant, Eq. (29) implies that ρ and therefore ρ is a constant. Then the second equation in DM DM (24) gives 4H λ2 0 ρ = , (30) DM κ2Q and therefore, from (29), we find 12H2λ2 ρ= 0 . (31) κ2Q2 Then by using the FRW equation (26), we find 4λ2 H = . (32) 0 3Q 1 4λ2 − Q2 (cid:16) (cid:17) This requires λ 1 < . (33) Q 2 By using (32), we can rewrite (30) and (31) as 16λ4 64λ6 ρ = , ρ= . (34) DM 3κ2Q2 1− 4Qλ22 3κ2Q4 1− 4Qλ22 2 (cid:16) (cid:17) (cid:16) (cid:17) Then we obtain Q2 1 4λ2 ρDM − Q2 = . (35) ρ (cid:16)4λ2 (cid:17) At the present time, ρ /ρ 1/3, and the fact that this ratio is of order unity today is called the coincidence DM ∼ problem. This observed ratio can be obtained in our model when λ2/Q2 3/16. ∼ De Sitter space can be realized by the a-independent energy density. The energy density of the phantom dark energy increases by the expansionbut it decreasesby the decay into the dark matter. On the other hand, the energy density of the dark matter decreases by the expansion but it increases by the decay of the dark energy. In the above solution, the decay of the dark energy into the dark matter balances with the expansion of the universe, and the energy densities of both the dark energy and dark matter become constant. This mechanism is essentially identical to one found in [8]. If the solution corresponding to de Sitter space-time is an attractor, the universe becomes asymptotic de Sitter space-time and any rip might be avoided. In order to investigateif the de Sitter space-time is an attractoror not, we consider the perturbation from the de Sitter solution in (32) and (34): 4λ2 16λ4 64λ6 H = +δH, ρ = +δρ , ρ= +δρ. (36) 3Q 1− 4Qλ22 DM 3κ2Q2 1− 4Qλ22 DM 3κ2Q4 1− 4Qλ22 2 (cid:16) (cid:17) (cid:16) (cid:17) (cid:16) (cid:17) Then the first FRW equation (26) gives 8λ2 δH =δρ+δ . (37) DM κ2Q 1 4λ2 − Q2 (cid:16) (cid:17) The conservation laws (24) and (25) give 16λ4 Q δρ˙ = δH δρ, κ2Q2 1 4λ2 − 2 − Q2 (cid:16)16λ4 (cid:17) 4λ2 δρ˙ = δH +Qδρ δρ . (38) DM DM −κ2Q2 1 4λ2 − Q 1 4λ2 − Q2 − Q2 (cid:16) (cid:17) (cid:16) (cid:17) 6 By eliminating δH in (38) using (37), we obtain Q 1 4λ2 2λ2 d δρ −2 − Q2 Q δρ dt δρDM = Q (cid:16)1 2λ2 (cid:17) 2Qλ2(cid:16)3−4Qλ22(cid:17)  δρDM . (39) (cid:18) (cid:19) − Q2 − 1−4λ2 (cid:18) (cid:19)  Q2   (cid:16) (cid:17)  In order for the de Sitter solution in (32) and (34) to be stable, all the eigenvalues of the matrix in (39) should be negative, which requires the trace of the matrix to be negative and the determinant to be positive, giving Q 1 4λ2 2Qλ2 3− 4Qλ22 <0, λ2 >0. (40) − 2 − Q2 − 1(cid:16) 4λ2 (cid:17) (cid:18) (cid:19) − Q2 The second condition can be trivially satisfied, and the first condition is also satisfied as long as (33) is satisfied. Therefore the de Sitter solution in (32) and (34) is stable and therefore an attractor. This tells us that the coupling of the dark matter with the dark energy as in (24) eliminates the little rip. Thus, if the universe where the dark energy dominates is realized, the universe will expand as in (8). If there is an interaction as given in (24), the dark energy decay into dark matter will yield asymptotic de Sitter space-time corresponding to Eq. (32). IV. SCALAR FIELD LITTLE RIP COSMOLOGY A. Minimally Coupled Phantom Models First consider a minimally coupled phantom field φ which obeys the equation of motion φ¨+3Hφ˙ V′(φ)=0, (41) − where the prime denotes the derivative with respect to φ. A field evolving according to equation (41) rolls uphill in the potential. In what follows, we assume a monotonically increasing potential V(φ). If this is not the case, then it is possible for the field to become trapped in a local maximum of the potential, resulting in asymptotic de Sitter evolution. Kujat, Scherrer, and Sen [16] derived the conditions on V(φ) to avoid a big rip, namely V′/V 0 as φ , and → →∞ V(φ) dφ . (42) V′(φ) →∞ Z p When these conditions are satisfied, w approaches 1 sufficiently rapidly that a big rip is avoided. − We now extendthis argumentto determine the conditions necessaryto avoida little rip. Clearly,we will haveρ → constantifV(φ)isboundedfromabove,sothatV(φ) V (whereV isaconstant)asφ . Wecanshowthatthis 0 0 → →∞ is alsoanecessarycondition. Suppose thatV(φ) is notboundedfromabove,sothat V(φ) asφ . Thenthe →∞ →∞ only wayfor the density ofthe scalarfieldto remainbounded isif the field “freezes” atsomefixed valueφ . However, 0 this is clearly impossible from equation (41), since it would require φ¨= φ˙ = 0 while V′(φ) = 0. Thus, boundedness 6 of the potential determines the boundary between little rip and asymptotic de Sitter evolution. Phantom scalar field models with bounded potentials have been discussed previously in Ref. [9]. B. Scalar-Tensor Models Using the formulation in Ref. [14], we now consider what kind of scalar-tensor model, with an action given by 1 1 S = d4x√ g R ω(φ)∂ φ∂µφ V(φ) , (43) − 2κ2 − 2 µ − Z (cid:26) (cid:27) can realize the evolutionof H given in Eqs. (8) or (11). Here ω(φ) and V(φ) are functions of the scalar field φ. Since the corresponding fluid is phantom with w < 1, the scalar field must be a ghost with a non-canonicalkinetic term. − If we consider the model where ω(φ) and V(φ) are given by a single function f(φ) as follows, 2 1 ω(φ)= f′′(φ), V(φ)= 3f′(φ)2+f′′(φ) , (44) −κ2 κ2 (cid:0) (cid:1) 7 the exact solution of the FRW equations has the following form: φ=t, H =f′(t). (45) Then for the model given by Eq. (8), we find 2λH 1 ω(φ)= 0eλφ, V(φ)= 3H2e2λφ+λH eλφ . (46) − κ2 κ2 0 0 (cid:0) (cid:1) Furthermore, if we redefine the scalar field φ to ϕ by 2eλ2φ 2H0 ϕ= , (47) κ λ r we find that the action (43) has the following form: 1 1 3λ2κ2 λ2 S = d4x√ g R+ ∂ ϕ∂µϕ ϕ4 ϕ2 . (48) − 2κ2 2 µ − 64 − 8 Z (cid:26) (cid:27) Note that in the action (48), H does not appear. This is because the shift of t in (8) effectively changes H . The 0 0 parameter A in [11] corresponds to 2λ/√3 in (8) and is bounded as 2.74 10−3Gyr−1 A 9.67 10−3Gyr−1, or 2.37 10−3Gyr−1 λ 8.37 10−3Gyr−1, by the results of the Supern×ova Cosmology≤Pro≤ject [15×]. In [11], it was × ≤ ≤ × shown that the model defined by Eq. (8) can give behavior of the distance modulus versus redshift almost identical to that of ΛCDM, so this model can be made consistent with observational data. As in Ref. [11], we can generalize the behavior of this model to H =H eCeλt. (49) 0 Here H , C, and λ are positive constants. Then we find 0 2 1 ω(φ)= H CλeCeλφeλφ, V(φ)= 3H2e2Ceλφ +H CλeCeλφeλφ . (50) −κ2 0 κ2 0 0 (cid:16) (cid:17) If we redefine the scalar field φ to ϕ by ϕ= √2H0Cλ dφeC2eλφeλ2φ = 1 8H0C eλ2φdxeC2x2 = 2 H0πErfi Ceλ2φ , (51) κ κ λ κ λ 2 Z r Z r "r # we may obtain the action where the kinetic term of the scalar field ϕ is +1∂ ϕ∂µϕ. In (51), Erfi[x] = Erf[ix]/i 2 µ with i2 = 1, where Erf[x] is the error function. In [11], it was shown that the model given by Eq. (49) can also be − consistent with the observations. As in Ref. [11], we can easily find models which show more complicated behavior of H such as H =H eC0eC1eC2eλt . (52) 0 On the other hand, in the model given by Eq. (11), we find ω(φ)= 2λH1e−λφ, V(φ)= 1 3 H H e−λφ 2+λH e−λφ , (53) − κ2 κ2 0− 1 1 n (cid:0) (cid:1) o and by the redefinition 2e−λ2φ 2H1 ϕ= , (54) κ λ r we find that the action (43) has the following form: 1 1 1 λκ2 2 λ2κ2 S = d4x√ g R+ ∂ ϕ∂µϕ 3 H ϕ2 + ϕ2 . (55) Z − (2κ2 2 µ − κ2 " (cid:18) 0− 8 (cid:19) 8 #) 8 In the action given by Eq. (55), H does not appear. This is because the shift of t in (11) effectively changes H . 1 1 Eq. (47) shows that in the infinite future t = φ + , ϕ also goes to infinity, that is, the scalar field climbs up → ∞ the potential to infinity. This climbing up the potential makes the Hubble rate grow and generates a rip due to the inertial force (1). On the other hand, Eq. (54) tells us that when φ + , ϕ vanishes. Note that the potential in → ∞ (55)isadoublewellpotentialsimilartothepotentialoftheHiggsfield,andϕ=0correspondstothe localmaximum ofthepotential. Therefore,inthe modelgivenby(55),thescalarfieldclimbsupthe potentialandarrivesatthelocal maximum after an infinite time. The behavior of the scalar field is different from that of the canonical scalar field, whichusuallyrollsdownthepotential. Thisphenomenonofhowthescalarfieldclimbsupthepotentialoccursdueto the non-canonical kinetic term. For the canonical scalar field ϕ , the field equation has the form of 2ϕ = V′(φ), but if the sign of the kinetic term is changed, we obtain 2ϕ =c V′(φ) for a non-canonical scalar fie∇ldt. cTha−t is, the ∇t c sign of the “force” is effectively changed. We now investigate the stability of the solution (45) in the model given by Eqs. (43) and (44) by considering the perturbation from the solution (45): φ=t+δφ(t), H =f′(t)+δh(t). (56) By using the FRW equations 3 1 1 1 H2 = ω(φ)φ˙2+V(φ), 2H˙ +3H2 = ω(φ)φ˙2 V(φ), (57) κ2 2 −κ2 2 − (cid:16) (cid:17) we find d δh 6f′(t) 6f′(t)f′′(t)+f′′′(t) δh dt(cid:18)δφ (cid:19)= −−3ff′′′((tt)) 3f′(t) !(cid:18)δφ(cid:19) . (58) In orderfor the solution(45) to be stable, allthe eigenvalues ofthe matrix in (58) shouldbe negative,which requires the trace of the matrix to be negative and the determinant to be positive, giving f′(t)f′′′(t) 3f′(t)<0, 3 >0. (59) − f′′(t) Thefirstconditionistriviallysatisfiedintheexpandinguniversesincef′(t)=H >0. Iftheuniverseisinthephantom phase, where f′′(t)=H˙ >0, the second condition reduces to f′′′(t)=H¨ >0. Then the model corresponding to (11) is unstable but the model corresponding to (8) is stable. There are no local maxima in the potential in (48), so one would expect the field to climb the potential well to infinity, generating a rip. In general,in a model which generates a big or little rip, H goes to infinity, which requires H¨ > 0. Therefore in the scalar field model generating a big or little rip, the solution corresponding to the rip is stable, and models that are asymptotically de Sitter can eventually evolve to have a rip. V. INCLUDING MATTER In the previous sections, we have neglected the contribution from matter except for the dark matter in Sec. III. In this section, we now consider the affect of additionalmatter components. We assume eachcomponent has a constant EoSparameterwi . Thentheenergydensityandpressurecontributedbyallofthesecomponentscanbeexpressed matter as ρmatter = ρi0a−3(1+wmiatter), pmatter = wiρi0a−3(1+wmiatter). (60) i i X X Here the ρi’s are constants. Even including these additional matter components, we can construct the scalar-tensor 0 model realizing the evolution of H by, instead of (44), ω(φ) = 2 g′′(φ) wmi atter+1ρia−3(1+wmiatter)e−3(1+wmiatter)g(φ), −κ2 − 2 0 0 i X V(φ) = 1 3g′(φ)2+g′′(φ) + wmi atter−1ρia−3(1+wmiatter)e−3(1+wmiatter)g(φ). (61) κ2 2 0 0 i (cid:0) (cid:1) X 9 Then the solution of the FRW equations (3) is given by φ=t, H =g′(t), a=a eg(t) . (62) 0 (cid:16) (cid:17) We may consider the example of (8), which gives a(t)=a0eHλ0eλt. (63) Then by using the FRW equations (3), we find the EoS parameter w corresponding to the dark energy is given by DE 3 H2 ρ w = κ2 − matter DE 1 2H˙ +3H2 p −κ2 − matter = (cid:16) κ32H02(cid:17)e2λt− iρi0a−03(1+wmiatter)e−3(1+wmiλatter)H0eλt . (64) −κ12 (2λH0eλt+3H02e2λPt)− iwmi atterρi0a0−3(1+wmiatter)e−3(1+wmiλatter)H0eλt When t becomes large, the contribution from the maPtter components decreases rapidly and w in (64) coincides DE with w in (9). The density parameter Ω of the dark energy is also given by DE ΩDE = κ32H23−Hρ2matter =1− 3κH22 ρi0a0−3(1+wmiatter)e−3(1+wmiλatter)H0eλt−2λt, (65) κ2 0 i X which rapidly goes to unity when t becomes large. It would be interesting to consider the cosmological perturbation in the model including the contribution from these matter components. Let t=0 represent the present. We now assume the matter consists only of dust with a vanishing EoS parameter. Then we find ρmatter = ρ0a−03e−3Hλ0 and pmatter = 0. Since ΩDE = 0.74, we find ρmatter = 0.26× 3κH202 by using (65). Since H is the Hubble parameter in the present universe, we find H = 7.24 10−2Gyr−1(≈70km/sMpc). Since 0 0 2.37 10−3Gyr−1 λ 8.37 10−3Gyr−1 (see below (48)), by using (64), w×e find 0.97 < w < 0.72, which DE could×be consistent≤with≤the ob×served value w = 0.972+0.061. − − DE − −0.060 VI. DISCUSSION Little rip models provide an evolution for the universe intermediate between asymptotic de Sitter expansion and models with a big rip singularity. We have shown that the EoS parameter w as a function of time is a less useful diagnostic of such behavior than is w as a function of the scale factor. As for the case of big rip singularities, a little rip can be avoided if the dark energy is coupled to the dark matter so that energy flows from the dark energy to the dark matter. Minimally coupled phantom scalar field models can lead to viable little rip cosmologies. The models we investigated that yield little rip evolution turned out to be stable against small perturbations, and we found that big rip evolution is also consistent with the conditions for stability. For phantom field models, rip-like behavior is an attractor. It is interesting that it was recently demonstrated that the little rip cosmology may be realized by a viscous fluid [17]. It turns out that the viscous little rip cosmology can also be stable. ScalarlittleripdarkenergyrepresentsanaturalalternativetotheΛCDMmodel,whichalsoleadstoanon-singular cosmology. It remains to consider the coupling of such a model with matter and to confront its predictions with observations. It is known [18] that in a local frame with a flat background, a classical field theory with w < 1 has a negative − kinetic energy term, and the corresponding quantum field theory has a tachyonic instability and a vacuum decay lifetime which appears finite, although possibly greater than the age of the universe. Our result shows that in the presence of a rip, the space-time expansion is so fast that this tachyonic instability does not have time to destabilize the global geometry and shows, interestingly, that the extraordinary conditions of a little rip can lead to an infinite lifetime. Acknowledgments P.H.F.andK.J.L.weresupportedinpartbytheDepartmentofEnergy(DE-FG02-05ER41418). S.N.wassupported inpartbyGlobalCOEProgramofNagoyaUniversity(G07)providedbythe MinistryofEducation, Culture, Sports, 10 Science & Technology (S.N.); the JSPS Grant-in-Aid for Scientific Research (S) # 22224003 and (C) # 23540296 (S.N.). S.D.O. was supported by MICINN (Spain) projects FIS2006-02842and FIS2010-15640,by CPAN Consolider IngenioProject,AGAUR 2009SGR-994,andJSPSVisitor Program(Japan)S11135. R.J.S.wassupportedinpartby the Department of Energy (DE-FG05-85ER40226). [1] S.Perlmutter et al. [SNCP Collaboration], Astrophys.J. 517, 565 (1999) [arXiv:astro-ph/9812133]; A.G. Riess et al. 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