The Cyclic Model of the Universe: A Theory of Regular Big Bangs and Dark Energy, Lecture notes of Physics

Download The Cyclic Model of the Universe: A Theory of Regular Big Bangs and Dark Energy and more Lecture notes Physics in PDF only on Docsity! The Cyclic Theory of the Universe Paul J. Steinhardt Department of Physics & Princeton Center for Theoretical Physics, Princeton University, Princeton, NJ 08544, U.S.A. Abstract The cyclic theory of the universe is a radical alternative to the standard big bang/inflationary scenario that offers a new approach for resolving the homogeneity, isotropy, and flatness problems and generating a nearly scale-invariant spectrum of fluctuations. The original formulation of the cyclic model was based on the picture suggested by M-theory in which the observable universe lies on a brane separated by a small gap along an extra dimension from a second brane. The cyclic model proposes that the big bang is a collision between branes that occurs at regular intervals; that each bang creates hot matter and radiation and triggers an epoch of expansion, cooling and structure formation; that there is an interbrane force responsible for drawing the branes together whose potential energy acts like dark energy when the branes are far apart; and that each cycle ends with the contraction of the extra dimension and a collision between branes – a new big bang – that initiates the next cycle. In more recent formulations, the cyclic model is realized with ordinary quantum field theory without introducing branes or extra dimensions. The key innovation common to all these models is the ekpyrotic phase, the period of ultra-slow contraction preceding the big bang. It is the ekpyrotic phase, rather than inflation, that is responsible for explaining the smoothness, flatness and large scale structure of the universe. It is also the ekpyrotic phase that generates the distinctive signatures in the spectrum of primordial gravitational waves and non-gaussian density fluctuations that will be used to test the cyclic model in forthcoming experiments. 1 Motivation The cyclic theory of the universe [Steinhardt and Turok(2002a), Steinhardt and Turok(2002b)] challenges the standard big bang/inflationary picture [Guth(1981), Linde(1982), Albrecht and Steinhardt(1982)] by offering an alternative explanation for the homogeneity, isotropy and flatness of the universe, the absence of magnetic monopoles and other massive stable relics from the Planckian era, as well as the generation of a nearly scale-invariant spectrum of density fluctuations. The theory is based on three underlying notions: (1) the big bang is not the beginning of space and time, but rather a transition from an earlier phase of evolution; (2) big bangs occurred periodically in the past and continue periodically into the future; and, (3) the key events that shaped the large scale structure of the universe occurred during a phase of slow contraction before the big bang. The third item contrasts with the standard big bang picture in which the large scale structure of the universe is shaped by a period of ultra-rapid expansion (inflation) that occurs strictly after the bang. The cyclic theory eschews inflation altogether in favor of an ekpyrotic phase of ultra-slow contraction before the bang during which w ≫ 1 (where w is the ratio of pressure to energy density). The big surprise and a compelling reason to take the cyclic theory seriously is the discovery that an inflationary expanding phase 1 with w ≈ −1 produces almost exactly the same large scale structure as an ekpyrotic contraction phase with w ≫ 1, to the degree that it is impossible to distinguish the two pictures based on current observations. Yet, the two pictures provide extraordinarily different visions of space, time, and the global structure of the universe, and are ultimately distinguishable through future observations. There are several important reasons for seeking an alternative to inflation. First, exploring alternatives is an effective way of determining if current observations are sufficient to specify uniquely the cosmic history of the universe. The example of the cyclic model teaches us that that answer is no. It also points us to the open issues theorists and experimentalists need to explore to move the field forward. Second, although the inflationary model was developed over two decades ago, its weaknesses have never been successfully addressed. Twenty years ago, the pressing questions were: What is the inflaton and why are its interactions finely-tuned? How did the universe begin and why did it lead to inflation? Those same ques- tions remain today. The hope had been that an improved approach to fundamental physics, such as string the- ory, might answer these questions. Yet, despite heroic efforts to construct stringy inflation models with tens or hundreds of moving parts (fluxes, throats and branes) and to consider a complex landscape of 10500 vacua, no compelling inflationary model has been found [Maldacena(2003)]. The most carefully studied examples to date (d-brane inflation in warped conifolds) suggest that inflation is typically spoiled by quantum corrections and topological constraints and that the few surviving examples require delicate tuning of both parameters and initial conditions [Baumann et al.(2007a)Baumann, Dymarsky, Klebanov, McAllister, and Steinhardt]. Other approaches to string inflation have not been critically examined to this degree, but there is no reason to expect that they will be immune. The discovery of dark energy is another setback for inflation. The hope had been that, by adding inflation to the big bang picture, the rest of the history of the universe would be set. We now know that, at best, inflation fixes the next 9 billion years. The rest of the history of the universe is dominated by dark energy, whose origin, stability and other physical properties are unknown. The standard big bang/inflationary picture appears to be put together piecemeal from separately adjustable components the big bang, inflation and, more recently, dark energy that are not directly linked to one another. This reduces the overall coherence and explanatory power of the theory. Perhaps the strongest motivation for considering alternatives to inflation is that the inflationary model may not have the powerful predictive power it was originally thought to have. The problem derives from the 2 Figure 1: A schematic illustration of the colliding brane picture of the cyclic theory beginning from the present epoch (marked “You are here). [Figure reprinted by permission of authors from Endless Universe, by P.J. Steinhardt and N. Turok, (Doubleday, 2007). important to appreciate, therefore, that most and perhaps all of the features of the colliding brane picture can be mimicked with ordinary quantum fields in three spatial dimensions (plus time). This “4d effective picture” of the cyclic model, using the same ingredients as the inflationary model but with different details, not only assures skeptics, but is a powerful tool for analyzing the quantitative behavior of the cyclic model. Hence, no matter what your point of view, both the 5d brane picture and the 4d effective picture are useful for analyzing the cyclic model. For this reason, we will present both in this overview. 2.1 The Colliding brane picture The cyclic theory of the universe is a direct outgrowth of an earlier cosmological model known as the “ekpy- rotic” model.[Khoury et al.(2001)Khoury, Ovrut, Steinhardt, and Turok, Khoury et al.(2002a)Khoury, Ovrut, Steinhardt, and Khoury et al.(2002b)Khoury, Ovrut, Seiberg, Steinhardt, and Turok] The ekpyrotic model was motivated by the idea that the collision between two brane worlds approaching one another along an extra dimension could create a hot big bang from a nearly vacuous space-time before the collision. The name, ekpyrotic, meaning “out of the fire, is drawn from ancient cosmology of the Stoic philosophers in which the universe is created from the ashes of a fire and, after a long period of cooling, is consumed by fire again. The contemporary 5 ekpyrotic concept is that the universe as we know it is made through a conflagration ignited by the collision of branes along a hidden fifth dimension. The ekpyrotic model is related conceptually to the pre-big bang model developed several years earlier [Gasperini and Veneziano(1993), Gasperini et al.(1997)Gasperini, Maggiore, and Veneziano, Gasperini and Veneziano(2003)]. Both share the idea that the big bang is not the beginning of space and time and that the density pertur- bations responsible for galaxy formation were generated before the bang. However, the pre-big bang model does not include an ekpyrotic contraction phase with w ≫ 1, which is crucial for smoothing and flatten- ing the universe and for generating scale-invariant perturbations, as detailed in Sec. 3). Consequently, the pre-big bang model has no smoothing mechanism and generates an unacceptably blue spectrum of density perturbations (unless other components are added). In this sense, the ekpyrotic scenario is a significant advance. The cyclic model differs from both the pre-big bang and ekpyrotic models because it is a comprehensive theory of the universe that incorporates the ekpyrotic theory of the big bang the dark energy phase observed today in a larger framework in which the two are intertwined. The big bang, instead of being a one-time event, repeats periodically every trillion years or so. This occurs through the regular collision of branes due to an interbrane force, plus the regular conversion of gravitational energy into brane kinetic energy and, ultimately, matter and radiation, as noted above. In the cyclic theory, the inter-brane potential energy density at present corresponds to the currently observed dark energy, providing roughly 70% of the critical density today. The dark energy and its associated cosmic acceleration play a role in restoring the Universe to a nearly vacuous state thereby allowing the cyclic solution to become an attractor. The attractor behavior means that the cycling is stable. Random deviations from perfect cycling due to quantum or thermal effects are dissipated by the smoothing and red shift effect of dark energy so that the universe returns to the ideal cycling solution after a few cycles. Note that dark energy, which is unnecessary and added ad hoc in the big bang/inflationary picture is here an essential element associated with the fundamental force that drives the entire cyclic evolution. In the colliding brane picture, the cyclic model assumes an overall geometry of two orbifolds planes, one with positive tension and one with negative tension, separated by an extra spatial dimension with negative cosmological constant, as suggested by heterotic M theory. The bulk space is warped, and the scale factor for each brane a± is purely a function of where the branes lie in the warped background. That is, when the 6 branes are apart, a± have different values, and, when they collide, a+ = a−. In the absence of matter or an interbrane force, the configuration is in static equilibrium. In the simplest treatment of the brane collision, the evolution is described solely by the time variation of the scale factors on the two branes, a±, where the ± distinguishes the orbifold planes with positive and negative surface tension. Then these two degrees of freedom can be translated directly into 4d effective picture using the mapping a+ = 2 a cosh((φ − φ∞)/ √ 6) a− = −2 a sinh((φ − φ∞)/ √ 6), (1) where a corresponds to the standard Robertson-Walker scale factor of the 4d effective theory and φ, known as the “radion,” is an ordinary 4d scalar field. These expressions are a dictionary for translating the 5d picture of colliding branes into a 4d effective picture of evolving quantum fields. The 4d radion field φ encodes information about the distance between branes in 5d: for example, the collision (a+ = a− in the 5d picture) translates into φ → −∞ in the 4d picture, and a finite gap corresponds to a finite value of φ. It is important to note that a(t0, the Friedmann- Robertson-Walker scale factor used to describe the expansion and contraction of the universe in the 4d effective picture, is neither the scale factor on positive tension brane (a+) nor the negative tension brane (a−), but rather a = 1 2 √ a2 + − a2 −. Consequently, even though the big crunch corresponds to a singularity (a → 0) in the 4d coordinates, the same event in 5d coordinates corresponds to non-singular, finite values of a±, whose evolution can be regular even when a vanishes. This feature is a sign of hope that the big bang singularity can translate into regular behavior in the 5d picture. We have noted that the branes move in a 5d space with negative cosmological constant, an anti-de Sitter or AdS space. For branes in AdS, because a+ and a− are the scale factors on the positive and negative tension branes, the coupling to matter in the action is a±ρ±m, where ρ±m is the matter density on the positive (negative) tension brane. In the 4d effective coordinates, the coupling induces a interaction between matter and the radion field proportional to β+ = 2cosh((φ − φ∞)/ √ 6 or β− = −2sinh((φ− φ∞)/ √ 6), respectively.. (The constant field shift φ∞ is arbitrary; it is convenient to choose φ = 0 to be the zero of the potential V (φ).) Since a+ and a− in low energy configurations are purely functions of their position in the AdS background and a− corresponds to smaller warp factor than a+, it is always the case that a− ≤ a+, so that a+ = a− is 7 Sec. 4). The scalar field φ satisfies φ̈ + 3Hφ̇ = −V,φ, (3) in the background (2), where dots denote derivatives with respect to t and H ≡ ȧ/a. Ignoring, for simplicity, the coupling between ordinary matter and φ, and the spatial curvature the Friedmann equation is H2 = 8πG 3 ( ρ + 1 2 φ̇2 + V (φ) ) (4) G is Newtons constant and ρ is the energy density of ordinary matter and radiation. The potential V (φ) is chosen by hand at present, but may be derivable from the higher dimensional theory [Steinhardt and Turok(2002a), Steinhardt and Turok(2002b), Buchbinder et al.(2007a)Buchbinder, Khoury, and Ovrut]. An example with the desired properties is V (φ) = V0 ( eαφ − e−βφ ) F (φ) (5) (see Fig. 3). Here V0 is equal to today’s dark energy density, α is non-negative (and typically ≪ 1) and β is positive (and typically ≫ 1). This ansatz for the V is motivated by string theory. When the branes are more than a few string lengths apart, the attractive force between them is expected to vary exponentially with the radion field φ with a coefficient that is itself φ-dependent. Here, for the purposes of illustration, we have approximated the φ-dependence in terms of two exponential terms in which the first dominates for large, positive φ and the second for negative φ. The factor F (φ) represents an effect that occurs when the branes are within a string lengthscale of one another. The extra dimension begins to disappear and a transition occurs in which the string coupling constant and all contributions to the potential fall rapidly to zero. For cosmology, the precise form of F (φ) is unimportant so long as it cuts off the steep exponential fall-off of the potential as φ moves from zero towards −∞. In the example in Fig. (3), the steep decline cuts off near a negative minimum, denoted Vend, where φ = φend. The moment that V bottoms out, even if it were to remain non-zero, it ceases to affect the cosmological evolution for the remaining time before the crunch. This is because the scalar field kinetic energy, which was increasing and remaining comparable to |V | during the ekpyrotic phase when V was steeply falling, continues to increase after V bottoms out. Consequently, the pressure and energy density both become kinetic energy dominated and their ratio w → 1 independent of the precise form of F . 10 B O U N C E φ V DARK ENERGY PHASE E K P Y R O T I C P H A S E EXPAND ING K E PHASE S ECOND EXPAND ING K E PHASE RAD IAT ION AND MAT TER PHASES V end CONTRAC T ING K E PHASE w ≈ 1 w ≈ 1 w ≈ 1 w = 1/3, 0 w ≈ –1 V 0 w > > 1 φ end 0 F I R S T B R I E F w>>1 PHASE Figure 3: An example potential V (φ). This plot shows where φ is on its potential at each stage in a cycle. The ratio of pressure to energy density is denoted by w. V0 is the potential energy density of the present dark energy phase. The future cycle evolution proceeds with φ moving from its present value (represent by the solid circle) towards the left, all the way to the bounce (φ → −∞) and the back to the right through a sequence of short-duration phases with w ≈ 1 or w ≫ 1, and finally settling back at the value of φ it has today, after which the cycle begins anew. Over the next nine billion years, the universe undergoes the usual radiation and matter dominates phases, followed by a dark energy dominated phase with w ≈= −1. This phase ends when field φ begins to roll slowly towards −∞ again and V becomes less than zero. Then begins the critical ekpyrotic contraction phase responsible for resolving the horizon and flatness problems and for generating a scale invariant spectrum of density perturbations; and it last until φ reaches φend at minimum of the potential V = Vend. After the ekpyrotic phase ends, the universe enters a scalar field kinetic energy dominated phase with w ≈ 1 that endures all the way to the bounce. 11 Of central importance to the cyclic model is the ekpyrotic phase, in which the universe is slowly con- tracting and the scalar field is rolling slowly down its steeply declining, negative potential. (In the original ekpyrotic model, this corresponds to the motion of a bulk brane towards a boundary brane; in the cyclic model, it corresponds to the motion of two boundary branes towards one another.) The equation of state for the universe is: w ≡ 1 2 φ̇2 − V 1 2 φ̇2 + V (6) which exceeds one when the potential energy density V is negative. For the sample potential, the ekpyrotic phase is the the period when the negative exponential dominates, V (φ) ≈ −V0e −βφ, and the background universe enters an attractor scaling solution, a(t) ∝ (−t)2/β2 ∝ eφ/β, H = 2 β2t ∝ −e−βφ/2, w ≈ β2/3 ≫ 1, (7) in which t is negative and increasing. Notice that, since β ≫ 1, the (4d effective) scale factor contracts very slowly compared to the Hubble radius |H |−1. So, the situation is the inverse of inflation during which the scale factor grows much more rapidly than the Hubble radius. The scaling solution is only relevant as long as φ > φend, and the function F (φ) is effectively unity. Once φ passes φend and heads towards −∞ (corresponding to the collision between branes), the potential energy is quickly converted to kinetic energy and the solution enters a kinetic energy dominated phase, with a(t) ∝ (−t) 1 3 ∝ eφ/ √ 6, H = 1 3t ∝ −e− √ 2/3φ, w ≈ 1. (8) This solution maintains itself all the way to φ → −∞. The scalar field reaches this boundary at infinity in finite time and rebounds. After the rebound, the solution followed is nearly the exact time-reverse of (8); the radiation and matter produced at the bang and a modest enhancement of the kinetic energy of φ have a negligible effect while φ < φend. There is a brief w ≫ 1 expanding phase right after φ passes φend moving to positive values, but the excess kinetic energy in φ quickly overwhelms the potential energy V (φ) and the universe enters a second expanding kinetic phase (Figure 1). The expanding w ≫ 1 phase is of modest duration and plays no significant role. Continuing into the expanding phase, the kinetic energy in φ redshifts away as a−6 and the universe becomes dominated by the radiation that was produced at the bounce. The net expansion in the entire kinetic 12 3.1 What is the w > 1 component? The effects due to a w > 1 energy component are critical to the success of the cyclic scenario. Given earlier attempts at oscillatory models over the last century, one might naturally wonder why the introduction of a w > 1 phase was not considered previously. The probable reason is that, prior to inflation, cosmologists often assumed for simplicity that the universe is composed of “perfect fluids” for which w = c2 s, where the equation of state w equals the ratio of pressure p to energy density ρ, and the speed of sound cs is defined by c2 s = dp/dρ. If w > 1 and the fluid is perfect, then cs > 1, which is physically disallowed for any known fluid. With the advent of inflation, cosmologists have become more sophisticated and flexible about what fluids they are willing to consider. The inflaton, for example, has w ≈ −1, yet the speed of sound is positive and well-behaved. A rolling scalar field with canonical kinetic energy has cs = 1. Similarly, it is possible to have w > 1 and yet 0 ≤ c2 s ≤ 1 without violating any known laws of physics. This opens the door to a novel kind of cyclic model. The w > 1 equation of state derives naturally the interbrane potential that draws the branes together or, equivalently, the effective potential V (φ) for the radion scalar field.. The interbrane separation is described by a modulus field φ with an attractive potential that is positive when the branes are far apart and becomes negative as the branes approach. The rolling from a positive to a negative value is necessary for switching the universe from accelerated expansion to contraction. To see how this occurs, consider the Hubble parameter after the universe is dominated by the scalar field and its potential: H2 = 8πG 3 [1 2 φ̇2 + V ] . (9) The universe is spatially flat after a period of accelerated expansion, so we have included only the scalar field kinetic and potential energy density terms. In order to reverse from expansion to contraction, there must be some time when H hits zero. Since the scalar field kinetic energy density is positive definite, the only way H can be zero is if V < 0. So, reversal from accelerated expansion forces us to have V roll from a positive value (where V as the dark energy) to a negative value. However, V < 0 immediately implies an equation of state w ≡ 1 2 φ̇2 − V 1 2 φ̇2 + V > 1. (10) An example is steep part of our sample potential in Eq. (5) and Fig. 3, V ∝ −e−βφ with β >> 1, which draws φ towards an attractor solution with constant equation of state w = (β2 − 3)/3 ≫ 1. Hence, the 15 w > 1 component was, from the beginning, an essential component of the ekpyrotic phase needed to turn the universe from expansion to contraction; and, fortunately, achieving this equation of state simple to achieve with scalar fields and simple potentials. The real miracle, though, is that the same feature has two other consequences that were unanticipated when it was first introduced and that are essential for making the cyclic model viable. 3.2 Solving the homogeneity, isotropy and flatness problems without inflation Before considering how an ekpyrotic contraction phase smooth and flatten the universe, it is useful to recall how an inflationary phase accomplishes the feat. In the standard big bang/inflation model, the universe is likely to emerge from the big bang with many ingredients contributing to the right hand side of the Friedmann equation: H2 = 8πG 3 [ρm a3 + ρr a4 + σ2 a6 + . . . + ρI ] − k a2 , (11) where H is the Hubble parameter; a is the scale factor; ρm,r is the matter and radiation density; σ2 measures the anisotropy; k is the spatial curvature and ρI is the energy density associated with the inflaton. The parameters ρi and σ are constants which characterize the condition when a = 1, which we can choose without loss of generality to be the beginning of inflation. Each energy density term decreases as 1/a3(1+w), for the value of w corresponding to that component. For the inflaton, we have assumed w ≈ −1 and the energy density is nearly a-independent. The “. . .” refers to other possible energy components, such as the energy associated with inhomogeneous, spatially varying fields. Inflation works because all other contributions, including the spatial curvature and anisotropy, are shrinking rapidly as a grows, while the inflaton density ρI is nearly constant. Once the inflaton dominates, the future evolution is determined by its behavior alone and its decay products. The result is a homogeneous, isotropic and spatially flat universe. Let’s consider the same equation in a contracting universe. The term that will naturally dominate is the one that grows the fastest as a shrinks. In this case, the anisotropy term, σ2/a6 appears to out. A more careful analysis including the full Einstein equation reveals that the universe not only becomes anisotropic, but also develops a large anisotropic spatial curvature. This triggers “chaotic mixmaster behavior,”[Belinskii et al.(1970)Belinskii, Khalatnikov, and Lifshitz, Belinskii et al.(1973)Belinskii, Khalatnikov, and Lifshitz, Demaret et al.(1986)Demaret, Hanquin, Henneaux, Spindel, and Taormina, Damour and Henneaux(2000)] re- sulting in unacceptably large inhomogeneity and anisotropy as the crunch approaches. 16 The story changes completely if there is an energy density component with w > 1 [Erickson et al.(2004)Erickson, Wesley, Steinhardt, The brane/scalar field kinetic energy density decreases as ρφ a3(1+w) , (12) where the exponent 3(1+w) > 6. Now, the scalar field density grows faster than the anisotropy or any other terms as the universe contracts. The longer the universe contracts, the more the scalar density dominates so that, by the bounce, the anisotropy and spatial curvature are completely negligible. Also negligible are spatial gradients of fields. The evolution is described by purely time-dependent factors, a situation referred to as ultralocal. The striking discovery is that a contracting universe with w > 1 has the same effect in homogenizing, isotropizing and flattening the universe as an expanding universe with w < −1/3. This point is worth empha- sizing, since some argued that the cyclic model relies on a dark energy phase to solve the horizon and flatness problems and, hence, should be regarded as just a variant of inflation [Kallosh et al.(2001)Kallosh, Kofman, and Linde, Linde(2002)]. It is now understood that dark energy is not needed for this purpose. In fact, if there were a period of dark energy expansion followed by a period of contraction with w < 1, the scenario would fail fail because, despite being rather homogeneous and flat at the end of the dark energy dominated phase, the universe would still enter the chaotic mixmaster phase before the crunch, leading to an unacceptable inhomogeneity, anisotropy and curvature. 3.3 w > 1 and scale-invariant perturbations As shown above, having w > 1 during the ekpyrotic contraction phase makes the universe homogeneous and isotropic, classically. The same condition is responsible for generating a nearly scale-invariant spectrum of density perturbations when the effects of quantum fluctuations are included [Khoury et al.(2003)Khoury, Steinhardt, and Turok Boyle et al.(2004)Boyle, Steinhardt, and Turok]. Here will treat the issue conceptually to make the compar- ison with inflation; a more detailed treatment is given in the next section. Both the cyclic model and inflation generate density perturbations from sub-horizon scale quantum fluctuations. In each picture, there is one phase when the sub-horizon scale fluctuations exit the horizon and a much later phase when they re-enter. In an exit phase with ǫ ≡ 3 2 (1 + w), the scale factor a(t) and the Hubble radius H−1 are related by the 17 like that of today’s universe, which drives the universe into a very low energy, homogeneous state. No new energy scale enters and the solution for the scalar field is then determined (up to a constant) by the scaling symmetry: φb = (2/β)ln(−At). Next, consider quantum fluctuations δφ in this background. The classical equations are time-translation invariant, so a spatially homogeneous time-delay is an allowed perturbation, φ = (2/β)log(−A(t+ δt)) → δφ ∝ t−1). On long wavelengths, for modes whose evolution is effectively frozen by causality, i.e. |kt| ≪ 1, we can expect the perturbations to follow this behavior. Hence, the quantum variance in the scalar field, 〈δφ2〉 ∝ ~t−2 (17) Restoring β via φ → βφ and ~ → β2 ~ leaves the result unchanged. However, since δφ has the same dimensions as t−1 in four spacetime dimensions, it follows that the constant of proportionality in (17) is dimensionless, and therefore that δφ must have a scale-invariant spectrum of spatial fluctuations. It is straightforward to check this in detail. Setting φ = φb(t) + δφ(t,x), to linear order in δφ the field equation is δ̈φ = −V,φφδφ + ∇2δφ, (18) and its solution as |kt| → 0 is scale invariant 〈δφ2〉 = ~ ∫ k2dk 4π2 1 k3t2 , (19) where the modes are labeled by wavenumber k. The generalization to two (or more) fields is straightforward. For example, consider two decoupled fields with a combined scalar potential Vtot = −V1e − R β1dφ1 − V2e − R β2dφ2 , (20) where β1 = β1(φ1), β2 = β2(φ2), and V1 and V2 are positive constants. For the purposes of illustration, we focus on scaling background solutions in which both fields simultaneously diverge to −∞. Then, as before, the fields each obtain nearly scale-invariant fluctuations; and, in particular, the relative fluctuation in the two fields, known as the “entropic perturbation δs ≡ (φ̇1δφ2 − φ̇2δφ1)/ √ φ̇2 1 + φ̇2 2 (21) satisfies [Gordon et al.(2001)Gordon, Wands, Bassett, and Maartens] δ̈s + ( k2 + Vss + 3θ̇2 ) δs = 0, (22) 20 where Vss = φ̇2 2 V,φ1φ1 − 2φ̇1φ̇2V,φ1φ2 + φ̇1 2 V,φ2φ2 φ̇1 2 + φ̇2 2 , θ̇ = φ̇2V,φ1 − φ̇1V,φ2 φ̇1 2 + φ̇2 2 . (23) Here, θ̇ measures the bending of the trajectory in scalar field space, (φ1, φ2). For the simplest case, a straight line trajectory with θ̇ = 0, the equation of motion (22) now becomes δ̈s + (k2 + V,φφ) δs = 0, (24) where we have set φ = φ1. This is exactly the same equation as that governing the fluctuations of a single field (18), and so, since δs is a canonically normalized field according to its definition in (21), the power spectrum of δs generated from quantum fluctuations is scale-invariant. Next we include gravity, so the action now becomes ∫ d4x √−g ( 1 16πG R − 1 2 N ∑ i=1 (∂φi) 2 − N ∑ i=1 Vi(φi) ) . (25) Then, the entropy perturbation equation (22) in flat spacetime is replaced (see e.g. [Gordon et al.(2001)Gordon, Wands, Bassett, by δ̈s + 3Hδ̇s + ( k2 a2 + Vss + 3θ̇2 ) δs = 4k2θ̇ a2 √ φ̇2 1 + φ̇2 2 Φ (26) where Φ is the Newtonian potential. For the straight trajectory θ̇ = 0, the entropy perturbation is not sourced by Φ and we can solve the equations rather simply; Ref. [Lehners et al.(2007a)Lehners, McFadden, Turok, and Steinhardt] for details. The result is a nearly scale-invariant spectrum with spectral index ns − 1 = 2 ǫ − ǫ,N ǫ2 . (27) where ǫ ≡ 3 2 (1 + w) ≡ φ̇1 2 + φ̇2 2 2H2 = β2 2(1 + γ2) , (28) and N ≡ ln(a/aend) is the running parameter that measures the number of e-folds remaining in the ekpyrotic phase. This epitomizes the spectrum of entropic perturbations generated during the ekpyrotic contraction phase. We next turn to how the entropic perturbations are converted into curvature perturbations after the ekpyrotic phase is complete. 21 4.2 Conversion of Entropic to Curvature Perturbations The generation of density perturbations in a contracting universe has been controversial due to a key differ- ence in the nature of growing and decaying modes in a contracting phase compared to an expanding phase. In general, there are two independent modes, curvature fluctuations (on comoving hypersurfaces) and time- delay perturbations. In inflation, the curvature fluctuation on comoving hypersurfaces is the growing mode, and the time delay is a decaying mode; also, fluctuations of the inflaton excite directly the curvature mode. However, the roles are reversed in a contracting phase: curvature fluctuations shrink to zero, and the time- delay modes grow as the bounce approaches; fluctuations of the radion excite directly the time-delay mode, but this does not, by itself, create fluctuations in temperature and density. Instead, a mechanism is needed to convert time-delay modes into curvature modes just before or at the bounce. The conversion is a significant issue. One of the theorems learned from studying perturbations in inflation is that the curvature fluctuation amplitude is conserved for modes outside the horizon [Bardeen(1980)]. If this were true for the cyclic model, then decaying curvature fluctuations before the bounce would imply negligible curvature fluctuations after the bounce (assuming no mode mixing at the bounce itself and the cyclic model would be inconsistent with observations. In particular, it can be shown that, if the cyclic picture could be reduced to a 4d effective field theory with a single scalar field (representing the radion), no method of conversion before the bound is known [Creminelli et al.(2005)Creminelli, Nicolis, and Zaldarriaga]. In the cyclic model, though, the 4d effective picture with one scalar field is only an approximation de- scribing a few of the degrees of freedom of the 5d colliding brane picture. The branes and the bulk separating them make the critical difference in the conversion process.[Tolley et al.(2004)Tolley, Turok, and Steinhardt, Craps and Ovrut(2004), Battefeld et al.(2004)Battefeld, Patil, and Brandenberger] First, the branes define a precise hypersurface for the bounce from big crunch to big bang – the time-slice in which each point on one brane is in contact with a point on the other brane [Tolley et al.(2004)Tolley, Turok, and Steinhardt, Steinhardt and Turok(2005)]. Since φ is the modulus field that determines the bounce between branes, one might imagine that this corresponds to a surface with uniform δφ = −∞, which is a comoving hypersurface. However, φ measures the distance between branes only in the case that they are static. If the branes are moving, there are corrections to the distance relation due the excitation of bulk modes that depend on the brane speed and the bulk curvature scale [Tolley et al.(2004)Tolley, Turok, and Steinhardt]. Roughly 22 considered above, the conversion occurs after the ekpyrotic phase has ended (t = tend) at time tref in the kinetic energy dominated phase when φ2 reflects off the boundary of moduli space. The resulting curvature perturbation spectrum is [Lehners and Steinhardt(2008a)] 〈R2〉 = ~ β2 1 |Vend| 3π2M2 Pl γ2 (1 + γ2)2 (1 + ln(tend/tref)) 2 ∫ dk k ≡ ∫ dk k ∆2 R(k) (31) for the perfectly scale-invariant case. Notice that the result depends only logarithmically on tref : the main dependence is on the minimum value of the effective potential Vend (see Fig. 3) and the parameter that determines the steepness of potential for φ1, β1. (The result depends on the steepness of the potentials for both φ1 and φ2, but the two are related by the condition φ̇2 = γφ̇1. Since γ = 1/ √ 3, β1,2 are of the same order of magnitude; henceforth, we will drop the subscript β ≡ β1.) Observations on the current Hubble horizon indicate ∆2 R(k) ≈ 2.2× 10−9. Ignoring the logarithm in (31), this requires β|Vmin| 1 2 ≈ 10−3MPl, or approximately the GUT scale. This is entirely consistent with heterotic M-theory. As for the spectral tilt, Eq. (27) becomes ns − 1 = 4(1 + γ2) β2M2 Pl − 4β,φ β2 , (32) where we have used the fact that β(φ) has the dimensions of inverse mass. The presence of MPl clearly indicates that the first term on the right is a gravitational term. It is also the piece that makes a blue contribution to the spectral tilt. The second term is the non-gravitational term and agrees precisely with the flat space-time result. For a pure exponential potential, which has β,φ = 0, the non-gravitational contribution is zero, and the spectrum is slightly blue, as our model-independent analysis suggested. For plausible values of β = 20 and γ = 1/2, say, the gravitational piece is about one percent and the spectral tilt is ns ≈ 1.01, also consistent with our earlier estimate. However, this case with β,φ precisely equal to zero is unrealistic. In the cyclic model, for example, the steepness of the potential must decrease as the field rolls downhill in order that the ekpyrotic phase comes to an end, which corresponds to β,φ > 0. If β(φ) changes from some initial value β̄ ≫ 1 to some value of order unity at the end of the ekpyrotic phase after φ changes by an amount ∆φ, then β,φ ∼ β̄/∆φ. When β is large, the non-gravitational term in Eq. (32) typically dominates and the spectral tilt is a few per cent towards the red. Our expression for the spectral tilt of the entropically induced curvature spectrum can also be expressed in 25 terms of the customary “fast-roll” parameters [Gratton et al.(2004)Gratton, Khoury, Steinhardt, and Turok] ǭ ≡ ( V MPlV,φ )2 = 1 β2 η̄ ≡ ( V V,φ ),φ. (33) Note that ǭ = 1/(2(1 + γ2)ǫ). Then, the spectral tilt is ns − 1 = 4(1 + γ2) M2 Pl ǭ − 4η̄. (34) By comparison, the spectral tilt for inflation is ns − 1 = −6ǫ + 2η (35) where the result is expressed in terms of the slow-roll parameters ǫ ≡ (1/2)(MPlV,φ/V )2 and η ≡ V,φφM2 Pl. Here we have revealed the factors of MPl to illustrate that both inflationary contributions are gravitational in origin. So, the range of spectral tilts for the simplest inflationary and ekpyrotic models are slightly different, with the ekpyrotic edging closer to ns = 1, but there is also considerable overlap, especially when more general potentials are considered. Finally, a distinctive prediction of ekpyrotic and cyclic models is that the density fluctuations are signif- icantly more non-gaussian than inflationary models. A general density fluctuation spectrum can be charac- terized by a series of n-point correlation functions 〈ρ(x1)ρ(x2) . . . ρ(xn)〉 (36) where 〈. . .〉 represents an average over all possible combinations of points x1 thru xn. The discussion above concerning scale-invariance and tilt referred specifically to the two-point function (also known as the power spectrum). If the spectrum is gaussian, all n-point fluctuations for odd n are zero and for even n are expressible as powers of the two-point function. Hence, all information in a gaussian spectrum is encoded in the two–point function. Both inflation and ekpyrotic/cyclic models predict a predominantly gaussian spectrum but also small non-gaussian contributions. These can be most easily detected by measuring the three-point function, also known as the bispectrum. And deviation from zero is a sign of non-gaussianity. In principle, the bispectrum, which measures a deviation from non-gaussianity known as skewness, can take many functional forms depending on the source of the non-gaussianity. In both the inflationary and ekpyrotic/cyclic models, the non-gaussianity is due to non-linear evolution of scalar fields that varies from point to point, so-called ‘local’ non-gaussianity. In this case, the non-gaussianity can be expressed as a cor- rection to the leading linear gaussian curvature perturbation, RL which, following the notational convention 26 used in [Langlois and Vernizzi(2007)], can be expressed as R = RL − 3 5fNLR2 L, where the sign convention for the coefficient fNL is the same as in Ref. [Komatsu and Spergel(2000)]. The parameter fNL is nearly scale-invariant and can have either sign in principle. A positive sign corresponds to positive skewness in the matter distribution (more structure) and negative skewness in the CMB temperature fluctuations (more cold spots). The magnitude of fNL is a measure of the degree of non-gaussianity. A key prediction of ekpyrotic/cyclic models is that fNL is O(100) or more times the value predicted by inflationary models (fNL . .1).∗ Although the precise value of fNL can vary significantly in both models, there is simple, intuitive rea- son why the non-gaussianity in ekpyrotic/cyclic models is generically several orders of magnitude greater than in inflationary models. The reason traces back to a defining feature that strongly distinguishes the two models: the difference in the equation of state during the period that density fluctuations are gen- erated. In standard versions of both models, the density perturbation spectra have their origin in scalar fields φi which develop nearly scale invariant perturbations while evolving along an effective potential V (φi). However, the potential is nearly constant during an inflationary phase in order to obtain winf ≈ −1 or, equivalently, εinf ≡ 3 2 (1 + winf ) ≪ 1; by contrast, the potential should be exponentially steep and negative to obtain εek ≫ 1, as required for an ekpyrotic phase. This means that the inflaton is nearly a free field with nearly gaussian quantum fluctuations. The non-gaussian amplitude depends on the deviation of the potential from perfect flatness or, equivalently, how close the slow-roll parameter εinf and its variation with time are to zero. This intuitive argument is consistent with the quantitative expression for fnLobtained for inflationary models [Maldacena(2003)]. However, a steep potential means that the scalar fields in the ekpyrotic model necessarily have significant nonlinear self-interactions whose magnitude depends just how large εek is. Because the magnitude of εek is O(100) or more times larger than εinf , the scalar field contri- bution to the non-gaussianity – which we will call the “intrinsic” part – is correspondingly larger for ekpy- rotic/cyclic models [Notari and Riotto(2002), Di Marco et al.(2003)Di Marco, Finelli, and Brandenberger, Buchbinder et al.(2007a)Buchbinder, Khoury, and Ovrut, Creminelli and Senatore(2007), Koyama et al.(2007)Koyama, Mizuno, Battefeld(2007), Lehners and Steinhardt(2008a)]. The magnitude of ε and its evolution during the last 60 e-folds of the inflationary or ekpyrotic/cyclic phase ∗N.B. This statement, which comes from the abstract of Ref. [Lehners and Steinhardt(2008b)], refers to the intrinsic contri- bution to fNL, as noted below, which is . .1 for inflation. Although the primary references by Lehners and the author state the condition correctly, a decimal point for the inflationary contribution was inadvertently missed in the earlier drafts of this review article. We regret any confusion this typo may have caused. 27 phase, both in terms of sign and magnitude. The error will improve modestly with further WMAP data, and then dramatically (σ ≤ 5) with the forthcoming Planck satellite mission and large scale structure studies. So, even though a precise value of fNL is not predicted, there is enough separation between the inflationary and ekpyrotic/cyclic predictions to make non-gaussianity a surprisingly good test. If observations of fNL lie in the range predicted by the intrinsic contribution of either inflationary or ekpyrotic/cyclic models, it is reasonable to apply Occam’s razor and Bayesian analysis to favor one cosmological model over the other. Combining with measurements of the spectral tilt significantly sharpens the test. The sign and magnitude of fNL provides additional information about the equation of state when the curvature perturbations were generated and, in some cases, the type of inflationary or ekpyrotic/cyclic model. To learn more about the detailed calculations behind the prediction, the reader is referred to Ref. [Lehners and Steinhardt(2008b)], which derives analytic estimates of the contributions to non-gaussianity. 5 Primordial Gravitational Waves in the Cyclic Model Although the spectra of density fluctuations are indistinguishable in ekpyrotic/cyclic and inflationary models to leading order, the spectra of gravitational waves to leading order are exponentially different. As a result, the search for primordial gravitational waves, either directly or their their imprint on the B-mode polarization of the cosmic microwave background radiation, is the most definitive way of distinguishing the two models. The difference arises because the gravitational backgrounds are so dissimilar in the ekpyrotic and infla- tionary phases when the fluctuations are generated. During inflation, the Hubble parameter H is large and nearly constant, resulting in strong gravitational Fluctuations and a nearly scale-invariant spectrum. During an ekpyrotic phase, H is roughly comparable to the current Hubble parameter H0, exponentially smaller than the case of inflation. Consequently, the gravitational effects are weak, and the primordial gravitational waves have an exponentially small amplitude. Also, H grows steadily throughout the ekpyrotic phase, which causes the gravitational wave spectrum to be blue, rather than scale-invariant. The method for computing the spectrum is detailed in Ref. [Boyle and Steinhardt(2005)]. Here we review the final results: The gravitational wave spectrum for the cyclic model can be divided into three regimes. There is a low frequency (LF) regime corresponding to long wavelength modes that re-enter after matter- radiation equality (or wavenumbers k < keq), and a medium frequency (MF) regime consisting of modes which re-enter between equality and the onset of radiation domination (keq < k < kr). (The dark energy 30 dominated phase has a negligible effect.) The spectrum for the dimensionless strain for these two regimes is: ∆h ≈ (kend/kr) 1 2 k2 0 πMplHα r { k−1+α (LF ) kα/keq, (MF ) (39) where k0 labels the mode whose wavelength is equal to the present Hubble radius, the subscript r refers to quantities evaluated at the beginning of the radiation-dominated phase after the bounce, and α ≡ 2/(β2 − 2) ≪ 1. The spectrum is blue with an amplitude that is exponentially small on horizon scales compared to the prediction for the simplest inflationary models. Finally, modes which exit the horizon during the ekpyrotic phase (before τend), and re-enter during the expanding kinetic phase (after the bound but before τr) result in a high frequency (HF) band (kr < k < kend): ∆h≈ (√ 2 π ) 3 2 (kendHr/kr) 1 2 −αk2 0 Mplkeqkr ∣ ∣ ∣ cos ( kτr− π 4 )∣ ∣ ∣ k 1 2 +α (HF ) (40) The spectrum can also be characterized in terms of Ωgw(k, τ0), the fractional energy density in primordial gravitational waves (ρgw)) per unit logarithmic wavenumber, in units of the critical density [Thorne(1987), Turner(1997)]: Ωgw(k, τ0) ≡ k ρcr dρgw dk = 1 6 ( k k0 )2 ∆h(k, τ0) 2 . (41) As shown in Fig. 5, Ωgw(k, τ0) in the cyclic model is very blue, with nearly all the gravitational wave energy concentrated at the high-frequency end of the distribution. By contrast, the primordial spectrum of gravitational waves in inflation is nearly flat. The calculation above shows that the primordial gravitational waves produced in the cyclic model are tens of orders below the detectable limit for a wide range of frequencies. An important correction to this predic- tion, though, is the indirect production of gravitational waves induced by the nearly scale-invariant spectrum of density perturbations.[Mollerach et al.(2004)Mollerach, Harari, and Matarrese, Ananda et al.(2007)Ananda, Clarkson, and Baumann et al.(2007b)Baumann, Steinhardt, Takahashi, and Ichiki] This effect is often ignored for infla- tionary models because it is small compared to the primordial gravitational wave contribution. How- ever, for the cyclic model, this “scalar-induced contribution is actually the leading source of gravita- tional waves at long wavelengths. A comparison of this scalar-induced contribution to the primordial spectrum for typical inflationary model (with a tensor-to-scalar ratio r = 0.1) is shown in Fig. 5. The scalar-induced contribution only applies to direct measurements of gravitational waves, such as BBO. The scalar-induced spectrum grows with time so that, extrapolating back to the decoupling of the cosmic back- ground radiation, the amplitude was too small to be detected. Hence, it left a completely negligible imprint 31 10 25 10 -10 10 -30 10 -50 10 -70 10 20 10 15 10 10 10 5 CMBPOL WMAP (2016 ?) BBO (2030 ?) INFLATION CYCLIC 10 0 10 5 10 10 10 15 Wavelength (meters) k (Mpc-1) 10 20 Figure 5: The fractional energy density in primordial gravitational waves per unit logarithmic wavelength in units of critical density for the cyclic and ekpyrotic models versus the wavelength of the gravitational wave. Superimposed are the current sensitivities for the Wilkinson Microwave Anisotropy Probe (WMAP) satellite and the sensitivities projected for future satellite missions to detect gravitational waves through their B- mode polarization effect on the cosmic microwave background (CMBPOL) and through direction detection in a space-based gravitational wave detector network such as the Big Bang Observatory (BBO). 32 regions of space are being created although any one region of space cycles for a finite time. Thanks to gravity, which provides an eternal source of energy and continuously creates more space, the cyclic model satisfies the conventional thermodynamic laws even if the cycling continues forever. Another suggestion has been that the holographic principle places a constraint on the duration of cycling [Steinhardt and Turok(2005)]. The argument is based on the fact that there is an average positive energy density per cycle. Averaging over many cycles, the cosmology can be viewed as an expanding de Sitter Universe. A de Sitter universe has a finite horizon with a maximal entropy within any observer’s causal patch given by the surface area of the horizon. Each bounce produces a finite entropy density or, equivalently, a finite total entropy within an observer’s horizon. Hence, the maximal entropy is reached after a finite number of bounces. (Quantitatively, a total entropy of 1090 is produced within an observer’s horizon each cycle, and the maximum entropy within the horizon is 10120, leading to a limit of 1030 bounces.) Closer examination reveals a loophole in this analysis [Steinhardt and Turok(2005)]. Although the overall causal structure of the four-dimensional effective theory may be de Sitter, it is punctuated by bounces in which the scale factor approaches zero. See Figure 7. Each bounce corresponds to a spatially flat caustic surface. All known entropy bounds used in the holographic principle do not apply to surfaces which cross caustics. Holographic bounds can be found for regions of space between a pair of caustics (i.e., within a single cycle), but there is no surface extending across two or more bounces for which a valid entropy bound applies. If the singular bounce is replaced by a non-singular bounce at a small but finite value of the scale factor, the same conclusion holds. In order for a contracting universe to bounce at a finite value of the scale factor, the null energy condition must be violated. However, the known entropy bounds require that the null energy condition be satisfied. Once again, we conclude that the entropy bounds cannot be extended across more than one cycle. Yet another way of approaching the issue is to note that both singular and non-singular bounces have the property that light rays focusing during the contracting phase defocus after the bounce, which violates a key condition required for entropy bounds. In particular, the light-sheet construction used in covariant entropy bounds [Bousso(2002)] are restricted to surfaces that are uniformly contracting, whereas the extension of a contracting light-sheet across a bounce turns into a volume with expanding area. Hence, if bounces are physically possible, entropy bounds do not place any restrictions on the number of bounces. Does this mean that the cycling model has no beginning? This issue is not settled at present. Fig. 7 suggests that the cyclic model has the causal structure of an expanding de Sitter space. For de Sitter space, 35 η χ 0 −π π Figure 7: The diagram shows the conformal time, η, which can vary between −π and zero, and the conformal spatial coordinate ξ, which varies between zero and π. The cyclic model has an average positive energy density per cycle, so its conformal diagram is similar to an expanding de Sitter space with constant density. The bounces occur along flat slices (curves) that, in this diagram, pile up near the diagonal and upper boundaries. For true de Sitter space, entropy bounds limit the total entropy in the entire spacetime. For the cyclic model, the bounds only limit the entropy between caustics (the bounces). A signal sent from the initial surface must travel through an infinite number of bounces to reach a current observer (dashed line). the expanding phase is geodesically incomplete. In principle, physical particles can travel along unobstructed paths through an empty de Sitter space from before the initial expanding surface to the present in finite proper time. For inflation, this picture is used to explain that the de Sitter expansion cannot have been ongoing eternally in the past and that information about conditions prior to inflation can plausibly make its way to a present-day observer. One cannot avoid the question of what happened “before inflation.” For the cyclic model, the story is perhaps different. We have already seen that periodic bounces change the entropy bounds on the cyclic model compared to pure de Sitter space. Perhaps the bounces need to be considered here, as well. The analogous paths backwards in time travel through an unbounded number of bounces, each of which is a caustic surface with a high density of matter and radiation. Any particle attempting this trip will be scattered and likely annihilated before reaching a present-day observer. Consequently, no information may pass from the asymptotic past (t → −∞) to the present, and what we see today may be totally, or almost totally, insensitive to the initial state. Even though Fig. 7 suggests that the cyclic model is geodesically incomplete, there is, perhaps, no physical meaning to “before cycling” if observers receive all information from the cycling epoch which contains an infinite number of cycles. This issue is currently unresolved. 36 7 What we have learned and what we need to learn What we have learned since the introduction of the cyclic theory of the universe six years ago has transformed it from a conceptual framework into a quantitative and predictive theory and has addressed most of the criticisms that have been raised. To appreciate the progress that has been made, it is worthwhile reviewing those criticisms and noting how the results described in the previous sections address them. • Does the cyclic model require fine-tuning of initial conditions? • Is the cyclic model subject to chaotic mixmaster oscillations as the big crunch approaches, which gen- erates unacceptable inhomogeneities and anisotropies? • Does the cyclic model rely on dark energy and accelerated expansion to resolve the horizon and flatness problems; and, if so, should it be regarded as simply a variant of inflation? The answer to all these question is a definitive no, as is clear from the discussion in Sec. 3. The ekpyrotic con- traction phase suffices to erase any inhomogeneity and anisotropy present as the dark energy dominated phase ends, even if the dark energy dominated phase lasts less than an e-fold [Erickson et al.(2004)Erickson, Wesley, Steinhardt, and T Chaotic mixmaster behavior is completely suppressed because w > 1. And since ekpyrotic contraction (and not dark energy) is entirely responsible for resolving the horizon and flatness problems and for generating the scale-invariant spectrum of density fluctuations, the model is clearly distinct from inflation. • Can the cyclic model generate a nearly scale invariant spectrum of density perturbations before the big bang? • Does the cyclic model predict any features in the density perturbation spectrum that differ from the simplest inflationary models? • Does the cyclic model generate an observable spectrum of gravitational waves? The answer to these questions is a definitive yes, as delineated in Secs. 4 and 5. The issue of generating curvature perturbations, controversial originally, has been settled with the introduction of the entropic mechanism discussed in Sec. 4, which treats the generation of perturbations with ordinary scalar fields in 4d using concepts familiar from inflationary cosmology. Non-gaussianity and scalar-induced gravitational waves are signatures of this mechanism and prime observational targets for testing the cyclic model. 37 The black holes are small and have a mass sufficiently small that the black holes decay in much less than one second, well before primordial nucleosynthesis. The black holes can be a boon to the scenario.[Baumann et al.(2007c)Baumann, Steinhardt, and Turok, Alexander and Meszaros(2007)] Their lifetime is long enough that they likely dominate the energy density before they decay. Consequently, their evaporation provides the entropy observed today. When they decay, their temperature rises near the end to values high enough to produce massive particles with baryon-number violating decays. At this point, the black holes are much hotter than the average temperature of the universe, so the decays occur when the universe is far from equilibrium. Assuming CP-violating interactions also exist, as in conventional high-temperature baryogenesis scenarios, the decay can produce the observed baryon asymmetry. The big bang/big crunch transition is the leading concern of cosmologists regarding the cyclic theory. If this hurdle can be convincingly surmounted, then the cyclic model will be generally recognized as a theory that not unifies cosmology in an appealing way, but also gives a new perspective on time, allowing the possibility that time exists long before the big bang and perhaps infinitely many cycles into the past. Note, though, that the cyclic model is compatible with their being a beginning. One could imagine the sudden creation from nothing of two infinitesimal spherical branes arranged concentrically, both of which undergo continuous expansion and collision. Both would grow enormously with every new cosmic cycle. After several cycles, the branes would appear very flat and parallel to any observer. Since all traces of previous cycles is diluted exponentially over the course of one period, there would be very little difference between this universe with a beginning and a universe with two branes colliding eternally. The emergence of the cyclic model as a viable competitor will also open minds to novel solutions to other theoretical issues, such as the cosmological constant (Λ) problem and the axion problem [Fox et al.(2004)Fox, Pierce, and Thomas Steinhardt and Turok(2006)]. For example, the cyclic model enables a new approach for solving the cosmo- logical constant problem in which Λ relaxes very slowly over course of many cycles from the Planck density to the exponentially smaller value it has today [Steinhardt and Turok(2006)]. The idea that Λ slowly relaxes has been tried within the context ofthe big bang/inflationary picture, but it is awkward to accommodate with inflation and dark energy. The relaxation must be very rapid during the first second of the universe to reduce Λ from the Planck scale to a level where it does not interfere with big bang cosmology, but this rapid reduction also tends to interfere with inflation; then, somehow, the relaxation must turn of at late times in 40 order to allow for the current period of accelerated expansion. Introducing an awkwardly tuned relaxation mechanism to resolve the cosmological constant fine-tuning problem does not seem like progress. The cyclic model introduces two important changes that free up constraints: (1) there is no inflation and, hence, no inflationary constraint on the relaxation rate; and, (2) the universe is much older, so it becomes to consider ultra-slow mechanisms for relaxing the cosmological constant with a characteristic timescale that is much greater then the Hubble time. is always far greater than the Hubble time. A possibility emerges in which the cosmological constant relaxes very slowly over the course of exponentially many cycles from a value near the Planck scale to the miniscule value observed today. One can then understand why the cosmological constant is exponentially smaller than one might guess based on the big bang picture; it is because the cosmological constant has had exponentially more time to relax than would be possible in the conventional big bang picture. A specific example of an ultra-slow dynamical relaxation mechanism is an idea first proposed by Larry Ab- bott [Abbott(1985)] in the context of standard big bang cosmology (see alsoRef. ([Brown and Teitelboim(1987)])) in which the vacuum energy density of a scalar field gradually decays through a sequence of exponentially slow quantum tunneling events, relaxing an initially large positive cosmological constant to a small value. The proposal introduced an axion-like scalar field with a “washboard” potential V (φ) = −M4 cos ( φ f ) + ǫ φ 2πf + Vother , (42) where M is the axion mass scale; f is of order the Planck mass; ǫ ≤ M4 is the coefficient of a linear term that breaks the shift symmetry (φ → φ + 2π); and Vother incorporates all other contributions to the cosmological constant that one seeks to cancel. It is natural for axions to have an exponentially M and ǫ because both terms in V are created by exponentially suppressed, non-perturbative physics. The potential has a set of equally spaced minima VN , with effective cosmological constants Λtotal = 8πGVN/3 spaced by 8πGǫ/3 (Fig. 8). No matter what value Vother has and no matter how large the initial vacuum density of the axion VN , the universe can tunnel its way down the washboard from minimum to minimum, canceling the vacuum density coming from all other sources, until the net cosmological constant reaches the smallest possible positive value Λtotal = 8πGV0/3. Abbott assumed the universe emerges from the big bang and quickly settles into a minimum with large positive VN , driving a period of inflation that dilutes any matter and radiation. Over time, the field φ 41 V-2 V -1 V 0 V 1 V 2 V 3 V 4 V 5 Figure 8: The effective cosmological constant for the washboard potential defined in Eq. 42 can take discrete values depending on which minimum VN the axion-like scalar field φ occupies. The difference in energy density between two successive minima, VN − VN−1 = ǫ, is exponentially smaller than the Planck density. 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