Inflation (cosmology) - Post‑Inflation Evolution

Reheating After Inflation The Purpose of Reheating After the inflationary period ends, the universe is not yet filled with the familiar particles we see today. Instead, most of the energy remains locked in the inflaton field—the quantum field whose dynamics drove inflation. Reheating is the crucial process that converts this inflaton field energy into a thermal bath of standard model particles, effectively launching the hot Big Bang phase. Think of reheating as inflation's finale: it transforms the inflaton's gravitational energy into heat and particles, creating the conditions necessary for the universe to evolve as we observe it today. Without reheating, there would be no radiation-dominated era and no pathway to the universe we inhabit. How Reheating Happens: Preheating The energy transfer from the inflaton to other fields doesn't occur gradually—it happens through an explosive mechanism called preheating. During preheating, the inflaton field oscillates around its minimum energy state. These oscillations trigger parametric resonance, a phenomenon where the oscillating inflaton couples to other quantum fields and causes them to grow exponentially. Parametric resonance is analogous to pushing a child on a swing at just the right time: each push (corresponding to each inflaton oscillation) adds energy, and the amplitude grows without bound. In the case of reheating, this resonance effect rapidly transfers energy from the inflaton to other particles—primarily to the "daughter" particles created by the coupling between fields. This exponential energy transfer happens far more efficiently than gradual decay would allow. The result is that within a very short time after inflation ends, most of the inflaton's energy has been channeled into a hot plasma of standard model particles. This marks the transition from the inflationary era to the radiation-dominated era. Constraints on Reheating Temperature Not all reheating scenarios are compatible with observations. The reheating temperature $T{\text{RH}}$ (the temperature of the particle bath immediately after preheating completes) must satisfy an important constraint: $$T{\text{RH}} \gtrsim 1 \text{ MeV}$$ This constraint arises from the success of primordial nucleosynthesis—the process by which the first light nuclei (primarily helium-4 and deuterium) formed in the early universe, roughly three minutes after the Big Bang. For nucleosynthesis to proceed with the abundances we observe today, the reheating temperature must be high enough that standard model particles are in thermal equilibrium during and after the nucleosynthesis epoch. If reheating were too cool, it would disrupt the delicate balance of nuclear reactions and produce the wrong elemental abundances. This observational constraint places a lower bound on how "cold" reheating can be, making it a crucial requirement for any inflationary model to match observations. Aftermath of Inflation: The Early Universe The Early Post-Inflation Epoch Once reheating completes (roughly one second after the Big Bang), the universe enters a new regime. It becomes dominated by radiation—ultrarelativistic particles like photons and neutrinos that move at speeds close to the speed of light. This is the radiation-dominated era. During radiation domination, the relationship between the scale factor $a(t)$ (which measures how much space has expanded) and cosmic time $t$ is: $$a(t) \propto t^{1/2}$$ This is fundamentally different from inflation, where the scale factor grew exponentially. The growth with the square root of time is much slower, and importantly, it represents decelerating expansion—the expansion rate $H = \dot{a}/a$ decreases with time. The universe is still expanding rapidly, but the expansion is gradually slowing down rather than speeding up as it did during inflation. This radiation-dominated phase lasted for roughly 50,000 years, until the universe cooled enough for protons and electrons to combine into neutral hydrogen (an event called recombination). <extrainfo> The square-root scaling $a(t) \propto t^{1/2}$ can be understood from the fact that in radiation domination, the energy density $\rho \propto a^{-4}$ (because both the number of particles per unit volume decreases as $a^{-3}$, and each particle's energy redshifts as $a^{-1}$). Combined with the Friedmann equation that relates expansion rate to density, this produces the $t^{1/2}$ scaling. </extrainfo> The Subsequent Evolution of the Universe The radiation-dominated era was not permanent. As the universe expanded and cooled, the energy density in matter (primarily in the form of protons, neutrons, and electrons) eventually exceeded the energy density in radiation. This marked the transition to matter domination. During matter domination (which lasted from roughly 50,000 years to 13 billion years after the Big Bang), the scale factor evolved as: $$a(t) \propto t^{2/3}$$ Matter domination was the era during which large-scale structure formed—gravity amplified tiny density fluctuations left over from inflation, eventually creating galaxies and galaxy clusters. However, observations from the 1990s onward revealed a dramatic surprise: the universe's expansion is currently accelerating, not decelerating. This acceleration is driven by dark energy, which now dominates the energy budget of the universe. We are currently in the dark-energy-dominated era, characterized by nearly exponential expansion. This stage will likely persist for the remainder of cosmic history. The sequence—inflation → reheating → radiation domination → matter domination → dark energy domination—represents the complete thermal history of the universe from the earliest moments to today.