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Story #3: Stellar evolution in 5 easy steps


Welcome back, gentle readers. After the first two instalments, in which I presented the beginnings of the Universe in a fashion that is hopefully understandable, I will move on to the next logical step – the stars. After all, stars are probably the most important objects in the Universe (of course, this depends on who you ask), and stellar evolution leads to three new topics that are fascinating and again very important. But I digress – everything will be covered in due course.


Step number one – Where do stars come from?


The answer to this question is simple. Stars appear as a result of gravitational contraction of gas clouds. Alright, this sounded more complicated than it is. The point is, there is a lot of hydrogen gas (with small amounts of other elements) floating around in the Universe. Before the first stars lit up, this was 99 percent hydrogen and no elements heavier than lithium; nowadays there are all elements of the periodic table (up to uranium) present in those clouds, although hydrogen still accounts for the vast majority of their mass. These clouds drifted around for quite some time and gradually cooled down until they reached temperatures of a few degrees Kelvin. At that point the thermal pressure (which arises from motion of warm particles) was so low that the clouds started collapsing under their own gravity. As they did so, their gravitational potential energy, which is negative and inversely proportional to the radius of the cloud, decreased. Total energy is, of course, conserved, – we assume that the cloud is isolated from other matter, – so this gravitational energy decrease results in increase of kinetic energy and release of radiation. In other words, the cloud starts rotating faster and shining. The rotation is also related to another fundamental conservation law – conservation of angular momentum. For a while this collapse is nearly isothermal – all random motions of the particles disappear as radiation. The constant low temperature also means that the Jeans mass of the cloud is very low – much lower than the solar mass – so the cloud will start to fragment. Typically, gas clouds are massive, several thousands of solar masses, if not more. In the isothermal collapse phase, they break up into smaller clouds of a solar mass or so. At this point, the cloud is already dense enough to become opaque to most radiation frequencies. Thus most of the radiation gets scattered inside the cloud and so the cloud heats up. Very little energy escapes the cloud – it is said to enter the adiabatic collapse phase (“adiabatic” is a sophisticated word for saying “not gaining or losing energy from the surroundings”). As the cloud’s temperature increases, its corresponding Jeans mass does so as well, until it becomes larger than the actual mass of the cloud. At this point the pressure inside the cloud is strong enough to balance the gravitational forces, thus stopping the collapse. The cloud becomes a proto-star and we reach step two.


Step number two – Protostars


This part of stellar evolution can be described with a large amount of mathematics or with a few sentences. I will go for the second approach, because the mathematics is rather lengthy and complicated, it is not exceptionally necessary and I do not remember the equations off the top of my head. Once the cloud reaches pressure-gravity balance – this state is called hydrostatic equilibrium and is central to explaining many processes and properties of stars – its temperature becomes rather constant. By that I mean that its luminosity scales as temperature to the 44th power, and this really is a lot. For a few thousand years the protostar continues to radiate energy and collapse isothermally, following the so-called Hayashi tracks on the H-R diagram*. Some protostars remain in this stage forever; they are called brown dwarfs and have masses ranging between 13 and 80 Jupiter masses. Most stars, however, reach a point when the internal pressure and temperature are large enough to allow hydrogen fusion to start. At this point, the star’s Main sequence starts.


Step number three – Main sequence


The vast majority of stars (not counting white dwarfs – more on them later) we know are in their main sequence. It is therefore concluded to be the longest part of the stellar lifetime. Stars in their main sequence may have masses ranging from less than 10 percent of the Sun’s up to hundreds of solar masses. The most common mass of a star is twice lighter than the Sun, but stars of our Sun’s mass are also rather common.


Main sequence stars all burn their internal hydrogen into helium. Depending on the actual process, they are grouped into Sun-like and massive (more than 1.5 solar masses) stars. Sun-like stars produce helium via a process called pp-chain: collision of two protons produces deuterium, which collides with another proton to produce helium-3 and finally another proton (or neutron) produces helium-4. Massive stars (above some 5 solar masses) burn matter via the CNO cycle: they produce some carbon early in their life, which then acts as a catalyst for protons to cling to, transforms into unstable isotopes of nitrogen and oxygen, which finally splits into a new carbon atom and a helium-4 atom.


Another property that distinguishes between the two kinds of stars is their internal structure. Sun-like stars have cores that are mostly transparent to the photons produced there (they are called radiative), while the outer layers are opaque (convective). In massive stars, these two regions are situated the other way around – the core is convective, while the envelope is radiative.


The more massive the star is, the less time it will spend on the main sequence. While our Sun has been here for the past 5 billion years and is expected to stay there for just as long, a star twice heavier will burn up ten times quicker and a star ten times more massive will only spend a few million years on the main sequence. Also, the star’s properties will slowly change while it is on the main sequence – it will contract and become hotter.


Once a star has finished burning hydrogen in its core, its physical parameters begin changing rather quickly, and it evolves off the main sequence.


Step number four – Post main sequence


There are many ways a star can evolve once it leaves main sequence. The two broad groups of evolutionary paths I will tell about here are representative of Sun-like and massive stars, respectively. Sun-like stars contract rapidly, thus increasing their temperature. Suddenly, the temperature becomes large enough for the star to burn helium-4 into carbon (the triple-alpha reaction). This reaction provides immense amounts of energy, quickly expanding the star to hundreds of times its original diameter. The star becomes a red giant. Once most of the helium is burnt – this takes a few million years, if that – the star’s core contracts again, but this time the contraction is more rapid, so the envelope does not contract and is lost. This lost envelope slowly spreads outward, becoming a planetary nebula (note that this has nothing to do with planets, it’s only a historical term). The core, now composed mainly of carbon, with some leftover helium and some oxygen, contracts to the smallest possible radius, which happens to be ten thousand kilometres or so. Such a remnant is called a white dwarf.


A massive star evolves similarly to the Sun-like one up until the helium burning ends. At this point, large stars have enough carbon so that it can start burning itself, as well as fusing with more helium. The star undergoes several flashes – the aforementioned helium flash, the beryllium flash, the carbon, oxygen, neon, magnesium flashes and so on, until either at the end of one flash the star is unable to move onto another, or the star fuses its core to iron. Iron is an interesting material in that its binding energy per nucleon (particle of the nucleus) is the largest, so it is a natural endpoint to both fusion and fission reaction chains. In this case, the star cannot produce any heavier elements and collapses onto itself. If it has not reached iron core yet, it usually collapses to a neutron star, whereas an iron-core star usually undergoes a supernova explosion, with the remnant being either a neutron star or a black hole.


Step number five – What is left


In the previous section, I briefly mentioned the three possible endpoints of stellar evolution – white dwarfs, neutron stars and black holes. The first two, if not in binary systems, slowly radiate away their energy and cool down. This process probably has to end somewhere, but it is so slow that the oldest white dwarfs have only cooled down to approximately 3000 K on the surface (the Sun’s surface temperature is 6000 K). At some point during the cool-down, white dwarfs crystallize and become huge diamonds (or other crystals, but carbon white dwarfs are the most abundant) floating in space. Black holes, on the other hand, float in space eating anything that comes too close. Well, not exactly, but we shall leave it at that.


Some more interesting things happen if these objects – called compact and degenerate** – are in binary systems. In this case, they can absorb matter from their companions – a process called accretion – and grow in mass. White dwarfs that grow beyond their theoretical mass limit, 1.4 solar masses, undergo a supernova explosion that leaves no remnant. Neutron stars may also produce supernovae or simply collapse to black holes. Black holes just grow in size until they might become supermassive black holes and harbour entire galaxies around them.


So this is it – a short overview of the main steps of stellar evolution. The subject is, as you may imagine, far more complicated than that, but this article hopefully sheds some light on our present understanding on the variety of stars we see in the Universe.


* – the HR, or Hertzsprung-Russell, diagram shows stellar temperature (sometimes called “colour”) on the horizontal axis and its luminosity (sometimes magnitude) on the vertical one.

** – degenerate in physics means several things, none of which is related to the colloquial meaning. In this case, degenerate matter is matter that has a significant part of its pressure arising not from thermal motion of particles, but rather from the Pauli Exclusion Principle as it applies to them.



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