KAHIBARO
Table of Contents
The first moments of the Universe
The early Universe refers to the time when the entire observable Universe was far hotter, denser, and more uniform than it is today. In this chapter, the main idea is that if the Universe is expanding now, then in the past everything must have been closer together. That means higher density and higher temperature. The early Universe is therefore studied as a sequence of stages in which temperature dropped as space expanded.
In very early times, matter as we know it did not exist in its familiar form. There were no atoms, no stars, no galaxies, and no planets. Instead, the Universe was filled with extremely energetic radiation and elementary particles interacting rapidly with one another. As expansion continued, the energy available in collisions decreased, and new structures gradually became possible.
A central rule of early-Universe cosmology is this: as the Universe expands, its average temperature decreases, and different kinds of particles and structures become possible at different temperature scales.
Why the early Universe was hot and dense
Imagine running cosmic history backward. Galaxies get closer together, the average density rises, and light becomes more energetic because the scale of the Universe becomes smaller. In this backward picture, the Universe approaches an extremely hot and compact state.
The early Universe is often described using a scale factor $a(t)$, which tells us how large the Universe is relative to some reference time. When $a(t)$ was much smaller than today, temperatures were much higher. A useful relation for a radiation-dominated Universe is
$$
T \propto \frac{1}{a(t)}
$$
This means that if the scale factor doubles, the temperature is roughly cut in half.
For much of the early Universe, smaller scale factor meant higher temperature:
$$
T \propto \frac{1}{a}
$$
This is one of the key links between expansion and thermal history.
The early Universe as a thermal system
In its earliest stages, the Universe behaved approximately like a thermal bath. Particles were constantly created, destroyed, and scattered. If interactions happened fast enough compared with the expansion rate, the contents stayed in thermal equilibrium. In equilibrium, particle energies followed statistical distributions set by temperature.
This picture is important because it lets physicists describe the early Universe using thermodynamics and particle physics together. Rather than tracking each particle individually, we describe the whole Universe by quantities such as temperature, density, and pressure.
At very high temperatures, radiation played the dominant role. In such conditions, the energy density of radiation scales strongly with temperature:
$$
\rho_{\text{rad}} \propto T^4
$$
So even a modest increase in temperature means a very large increase in energy density.
A timeline of major early stages
The early Universe is usually divided into eras. Exact details at the very earliest times remain uncertain, but the broad thermal history is well established for later stages.
| Approximate time after beginning of expansion | Typical condition | Main feature |
|---|---|---|
| Less than $10^{-43}\,\text{s}$ | Unknown quantum gravity regime | Present theories incomplete |
| Around $10^{-36}\,\text{s}$ or earlier | Extremely high energy | Possible inflationary era |
| $10^{-12}\,\text{s}$ | Very hot particle plasma | Electroweak processes important |
| $10^{-6}\,\text{s}$ | Quarks confined into hadrons | Protons and neutrons begin to form |
| About $1\,\text{s}$ | Neutrinos decouple | Weak interactions become less effective |
| From a few seconds to a few minutes | Cooling nuclear plasma | Light nuclei begin to form |
| About $380{,}000\,\text{years}$ | Much cooler ionized gas | Atoms later form, linked to the CMB |
This chapter focuses on the earliest thermal phases themselves, especially before atoms and stars existed.
The Planck era
The earliest interval is called the Planck era, roughly before
$$
t \sim 10^{-43}\,\text{s}
$$
At such times, temperatures and energies were so extreme that both quantum physics and gravity would have been equally important. Our current theories, general relativity and quantum mechanics, are not yet fully unified in this regime. Because of this, we do not have a complete, experimentally confirmed description of the Universe at that stage.
So the Planck era marks a boundary of present understanding. It is not that nothing can be said, but rather that our usual equations are not reliable enough there.
Before about $10^{-43}\,\text{s}$, known physics is incomplete. This interval is called the Planck era.
Inflation
Many cosmological models include a very brief period called inflation. During inflation, the Universe expanded extraordinarily rapidly in a tiny fraction of a second. This was not ordinary motion through space. It was space itself stretching.
Inflation is important because it helps explain several striking features of the Universe. It can explain why the Universe appears so nearly flat, why distant regions have nearly the same temperature, and why large-scale structure could grow from tiny initial irregularities.
A simple qualitative picture is that a small patch of space expanded enormously, smoothing out curvature and spreading tiny quantum fluctuations to cosmic scales. Those tiny fluctuations later became the seeds of galaxies and clusters of galaxies.
The exact mechanism of inflation belongs to advanced cosmology and particle physics, but for beginners it is enough to understand that inflation is a proposed ultra-rapid expansion stage in the very early Universe.
The particle soup
After the earliest unknown stages, the Universe was filled with a hot plasma of elementary particles. At sufficiently high temperatures, collisions were energetic enough to create particle-antiparticle pairs. Matter and radiation were tightly coupled.
Instead of atoms, there were quarks, leptons, photons, neutrinos, and other high-energy particles. Because temperatures were so large, rest mass often mattered less than thermal energy. A useful comparison is between thermal energy and particle rest energy:
$$
k_B T \sim mc^2
$$
When $k_B T$ is comparable to or greater than $mc^2$, a particle species can be produced abundantly in the thermal bath.
As the Universe cooled, heavier particles became harder to create. Many particle-antiparticle pairs annihilated, leaving behind radiation and the small excess of matter that survived.
A particle species is abundant in the early thermal bath when the thermal energy is high enough:
$$
k_B T \gtrsim mc^2
$$
Cooling removes the ability to create the heaviest particles first.
Quarks, hadrons, and the formation of protons and neutrons
At very early times, quarks moved in a dense plasma. As temperature dropped, quarks became confined into hadrons. The most important stable hadrons for later cosmic history are protons and neutrons.
This transition happened when the Universe was still extremely young, roughly around microsecond timescales. After that, free quarks were no longer the normal state of matter on large scales. Instead, the Universe contained protons, neutrons, electrons, neutrinos, photons, and their antiparticles where conditions allowed.
The relative numbers of protons and neutrons became especially important, because these would later determine how much helium and other light nuclei could form.
Freeze-out
A very important idea in early-Universe physics is freeze-out. While particle interactions are fast, thermal equilibrium is maintained. But the Universe is expanding, and expansion reduces both density and temperature. Eventually, some interaction rates become too slow to keep up.
If the interaction rate is called $\Gamma$ and the expansion rate is called $H$, then equilibrium is maintained roughly when
$$
\Gamma \gg H
$$
and freeze-out occurs when
$$
\Gamma \lesssim H
$$
At that point, a particle species or a process stops tracking equilibrium closely.
One important example is weak-interaction freeze-out, which affects the neutron-to-proton ratio. Another is neutrino decoupling, when neutrinos stop interacting frequently with other particles and begin to travel almost freely through space.
Freeze-out condition:
$$
\Gamma \lesssim H
$$
When interactions become slower than cosmic expansion, equilibrium can no longer be maintained.
Neutrino decoupling
When the Universe was about one second old, neutrinos interacted too weakly and too infrequently to remain tightly coupled to the rest of the plasma. They effectively decoupled and began streaming freely through the Universe.
This produced a background of relic neutrinos, somewhat like the cosmic microwave background but much harder to detect. These neutrinos still exist today in principle, though direct observation is extremely difficult.
Neutrino decoupling is one of the important milestones of the early Universe because it shows how different particle species stopped sharing energy at different times.
Matter and antimatter
The early Universe likely contained both matter and antimatter in enormous quantities. When a particle meets its antiparticle, they can annihilate into radiation. If matter and antimatter had been created in exactly equal amounts and remained perfectly symmetric, almost all of both would have annihilated away, leaving a Universe with very little ordinary matter.
But the visible Universe clearly contains matter. This means there must have been a small excess of matter over antimatter in the early Universe. After most annihilations occurred, that tiny excess remained and became the matter that later formed stars, galaxies, and people.
The origin of this asymmetry is still an area of active research. For beginners, the key point is simple: ordinary matter survived because the early Universe was not perfectly symmetric between matter and antimatter.
Radiation domination
In the earliest well-understood stages, radiation contributed more to the total energy density than matter did. This is called the radiation-dominated era. In this regime, the expansion behavior differs from that in a matter-dominated Universe.
During radiation domination, the scale factor evolves approximately as
$$
a(t) \propto t^{1/2}
$$
This relation is one reason the first moments passed through dramatic changes very quickly. A small increase in time could correspond to a large drop in temperature.
In a radiation-dominated Universe,
$$
a(t) \propto t^{1/2}
$$
and the temperature falls rapidly as expansion proceeds.
The first nuclei
As the Universe cooled further, conditions finally became suitable for protons and neutrons to combine into light nuclei. This happened within the first few minutes. The detailed study of those abundances belongs to primordial nucleosynthesis, but it is important here to see why this event belongs naturally to early-Universe history.
Before that moment, the temperature was so high that nuclei could not survive for long. High-energy photons would quickly break them apart. Only after enough cooling did stable light nuclei become possible.
This shows a general rule of the early Universe. Expansion cools the cosmos, and cooling opens the door to new forms of structure.
What the early Universe was not
It is common to imagine the early Universe as an explosion into empty space. That picture is misleading. The Big Bang model does not describe matter exploding outward into a pre-existing void. Instead, it describes space itself expanding everywhere.
Every region was once denser and hotter. There was no special center of the expansion within the observable Universe. The increase in distances occurred throughout space.
Limits of our knowledge
The early Universe is partly well tested and partly uncertain. Once the Universe cools to stages connected with light-element formation and later observable relics, our confidence becomes much stronger. At still earlier times, especially near the Planck era, knowledge becomes more speculative.
This means cosmology combines established physics with open questions. We know the Universe was once hot and dense. We have strong evidence for a thermal history and for expansion. But some of the very earliest mechanisms, such as the detailed physics of inflation or the exact origin of matter asymmetry, are not yet fully known.
The essential picture
The early Universe was a hot, dense, rapidly expanding physical system. As it expanded, it cooled. As it cooled, different particles stopped interacting in equilibrium, some species disappeared, others survived, and the first simple nuclear structures became possible. This sequence laid the foundation for everything that came later.
Essential summary of the early Universe:
$$
\text{Expansion} \Rightarrow \text{Cooling} \Rightarrow \text{Changing particle physics} \Rightarrow \text{Formation of new structures}
$$
This chain of events is the core idea of early-Universe cosmology.