by Frank Close
To find the answer it was necessary to make billions of Bo and , and to study them in detail. To do so, ‘B-factories’ – accelerators of e– and e+ that collide at a total energy of around 10 GeV where Bo and are copiously produced – were designed and built in California and in Japan. These are relatively compact machines on the scale of modern particle physics, being only a few hundred metres in circumference, but involving high-intensity beams of current controlled with greater precision than ever achieved before.
The accelerators were completed in 1999 and after initial testing began to collect data. To get definitive results requires creating and studying vast numbers of the bottom particles. It is like tossing a coin: chance might make it come up heads five or even ten times in a row, but if this continues to happen, then something is special about the coin. So it is with the study of ephemeral subatomic particles. They live for less than the blink of an eye and it is what remains after they die, their fossil relics if you like, that have to be decoded. One needs to have huge numbers of such fossils in order to tell if any differences are real or the result of chance.
There are many varieties of fossils that can be studied, and specialist teams at the two accelerators have begun to collect and measure the characteristics of several of these. Among them is a particular species, known as the ‘psi-K-short’ events – where Bo or decay and leave ψ and a particular mix of Ko and – that theorists predicted would be the most immediate indicator of a difference between bottom matter and bottom antimatter. By 2003 it was clear that these decays do show a large difference between matter and antimatter, as predicted. It will take several years of studying the properties of bottom particles to establish whether they hold the full answer to the matter-antimatter conundrum for the large-scale asymmetry in the basic seeds of matter, or whether the asymmetry exhibited by the strange and bottom particles is just an arcane phenomenon in the exotic forms of particles.
Chapter 9
Where has matter come from?
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Stars are cookers of heavy elements out of raw hydrogen. The seeds of hydrogen are in the quarks. What do we know about the behaviour of particles in the early universe?
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We exist because of a series of fortunate accidents: the fact that the Sun burns at just the right rate (faster and it would have burned out before intelligent life had the chance to develop; slower and there might not have been enough energy for biochemistry and any life at all); the fact that protons – the seeds of hydrogen – are stable, which enables stars to cook the chemical elements essential for the Earth to be built; and the fact that neutrons are slightly heavier than protons, which enables beta radioactivity, transmutation of the elements such as the protons of hydrogen into helium, which in turn enables the Sun to shine. Were any of these, or several others, slightly changed, we would not be here.
We and everything are made from atoms. Where did these atoms come from? Most recently (by which I mean 5 billion years!) they were formed inside a long-dead star where they were all cooked from protons, the nuclei of the simplest atomic element – hydrogen. The protons were formed very early in the universe and its constituent quarks, and also the electrons, were made within the first moments. This chapter describes how the stuff that we are made of came to be.
It is primarily protons that form the Sun and fuel it today. Let’s first describe how the Sun works and provides the energy for us to exist.
Hydrogen is the simplest atom, where a single negatively charged electron encircles a central positive proton. Hydrogen may be relatively uncommon on Earth (except when trapped inside molecules such as water – H2O) but in the universe at large it is the most common atomic element of all. At Earthly temperatures atoms can survive but at higher temperatures, above a few thousand degrees, the electrons are no longer trapped but roam free: the atom is said to be ionized. This is what it is like inside the Sun: electrons and protons swarm independently in the state of matter known as plasma.
Protons can bump into one another and initiate a set of nuclear processes that eventually converts four of them into the nuclei of the next simplest element: helium. The energy locked into a single nucleus of helium (its E = mc2) is less than that in the original four protons. This ‘spare’ energy is released into the surroundings, some of it eventually providing warmth here on Earth.
The protons have to touch in order to fuse and build up helium. This is hard as their positive charges tend to repel them, keeping them apart. However, the temperature of 10 million degrees gives them enough kinetic energy that they manage to encroach near enough to start the fusion power process. But it is only just enough: 5 billion years after its birth, any individual proton has only a 50:50 chance of having taken part in the fusion. Put another way: this far the Sun has used up half of its fuel.
This is the first fortunate circumstance. Humans are the pinnacle of evolution and it has taken almost all of those 5 billion years for us to emerge. Had the Sun burned faster, it would have died before we arrived.
So let’s see what happens and then why it is balanced just right.
The first step is when two protons meet and touch. One of them undergoes a form of radioactive decay, turning into a neutron and emitting a positron (the antiparticle of an electron) and a neutrino. Normally it is the neutron that decays, due to its extra mass and associated instability, into a proton, electron, and neutrino. An isolated proton being the lightest baryon, by contrast, is stable. But when two protons encroach, they feel electrostatic repulsion; this contributes to their total energy making it exceed that of a deuteron (a proton and neutron bound together). As a result one of the protons can turn into a neutron, which then binds to another proton, increasing the stability. This decay of the proton leads to a neutron, neutrino, and positron, the positively charged antiparticle of an electron.
So the very first part of the solar fusion cycle produces antimatter! The positron is almost immediately destroyed as it collides with an electron in the plasma, producing two photons which are scattered by the electrically charged plasma, eventually working their way to the solar surface (this takes several thousand years), by which time their energy is much reduced and they help form part of sunlight. The neutrinos pour out from the centre unhindered and reach us within a few minutes.
So what has become of the neutron and proton? They grip one another tightly, courtesy of the strong nuclear force, and bind together: this doublet is a nucleus of heavy hydrogen – the deuteron. This deuteron finds itself in the midst of a vast number of protons, which still form the bulk of the Sun. Very rapidly the deuteron links with another proton to make a nucleus of helium: helium-3. Two of these helium-3 can join and rearrange their pieces to form a nucleus of helium-4 (the stable common form), releasing two spare protons.
So the net result of all this is that four protons have produced a single helium, two positrons, and two neutrinos. Protons are the fuel, helium the ash, and the energy is released in the form of gamma rays, positrons, and neutrinos.
30. Converting hydrogen to helium in the Sun.
The latter steps, where a deuteron and a proton make 3He and then lead to 4He happen almost instantaneously; it is the tardiness of the first step, p + p → dνe+ that controls the (slow) burning of the Sun that has been so important for us.
The rate of the burning depends upon the strength of the weak force, which transmutes the proton into a neutron (‘inverse beta decay’). This force has parallels with the electromagnetic force, as described earlier. The electromagnetic force is transmitted by photons, which are exchanged between one electrically charged particle and another. Photons are massless: this enables them to spread to large distances without restrictions from energy conservation and hence gives the electromagnetic force a long range. The weak force, by contrast, owes its feebleness (at least at the energies characteristic of Earth and the Sun) to the large mass of the W boson and its consequent restricted range.
So the slowness of solar burning is cont
rolled by the feebleness of the weak force which is in turn controlled by the large mass of the W boson. Had its mass been smaller, the effective strength of the ‘weak’ force would have been stronger and rate of solar burning faster. Why does the W mass have this fortunate value? We do not know. We do not even know for sure where mass actually comes from though there are ideas due to Peter Higgs that will be tested very soon (in Chapter 10).
There are other examples where masses play a sensitive role in determining our fate. As we have discussed above, beta decay involves a neutron turning into a proton and emitting electron and neutrino. This requires the neutron to be heavier than the proton – which it is, whereby protons are the stable seeds of atoms and chemistry. (Had neutrons been lighter then it would have been neutrons that emerged as the stable pieces from the Big Bang. These neutral particles would have been unable to attract electrons to form atoms so chemistry would have been different or non-existent.) The neutron is only one part in a thousand heavier than the proton, but this fortunately is enough that an electron can be produced, or put another way, the electron mass is small enough that it can be produced in such a process. Had it been larger then beta decay and the Sun would have been frozen; had it been smaller, beta decay would have been faster, the Sun’s dynamics different, the intensity of ultraviolet light higher and unhealthy for us. (The mass of the electron helps determine the size of atoms such as hydrogen; smaller mass correlates with a larger atom and vice versa. So things have the size they do in part because the mass of the electron is as it is.) The reason for this pattern of masses is still to be found.
So the Sun is shining courtesy of nuclear fusion. In another 5 billion years its hydrogen will all have gone, turned into helium. Some of the helium already is itself fusing with protons and other helium nuclei to build up the nuclear seeds of heavier elements. These processes also produce neutrinos, some of higher energies than those produced in the primary proton fusion; and so by detecting neutrinos from the Sun, and measuring their energy spectrum, we can begin to get a quantitative look inside our nearest star.
Five billion years hence these will be the primary processes, along with the fusion to build up yet heavier elements. In some stars (but not our Sun) this process continues, building up the nuclei of elements up to iron, which is the most stable of all (there are even elements beyond iron that are built but they tend to be rarer). Eventually such a star is unable to resist its own weight, and it collapses catastrophically. The shock waves spew out matter and radiation into space. This is known as a supernova. So stars begin as hydrogen, and with these ingredients they cook the periodic table; a supernova is the agent that pollutes the cosmos with the nuclear seeds of these chemicals.
So where did the material for the primary stars come from?
The early universe
The basic pieces of nuclear matter, quarks, emerged from the big bang along with electrons. The universe rapidly cooled so that the quarks clustered together to form protons. The following processes took place:
e(electron) + p(proton) n (neutron) + ν(neutrino)
The double arrow is to illustrate that this process could occur in either direction. The neutron is slightly heavier than the combined masses of a proton and an electron, so the ‘natural’ direction for the processes was to go from right to left: the neutron has a natural tendency to lower the mass of the whole, liberating energy via E = mc2. However, the heat of the universe was such that the electrons and protons had considerable amounts of kinetic energy such that their total energy exceeded that locked into the mass (mc2) of a neutron. So in these hot conditions the process could as easily run from left to right (electron and proton converting into neutron and neutrino) as the other direction where the neutrons and neutrinos turned back into their electrically charged cousins. In these circumstances we say that the universe was in thermal equilibrium.
But the universe was rapidly cooling, which made it harder for the production of neutrons to continue. After a microsecond the universe had cooled to a point where this neutron production reaction was effectively frozen out. The surviving reaction was
During this epoch any neutrons that had been produced in the earlier heat would be dying out. Every ten minutes their numbers halved (we say they have a ‘half life’ of about ten minutes). There is no longer enough energy to replace them. But not all the neutrons died as some fortunate ones bumped into protons, whereupon they fused to one another to make a deuteron (a bound system of a single proton and a neutron which is lighter than an isolated proton and neutron are).
At this stage the universe at large plays out the sequence that is going on in the sun today: deuterons and protons building up nuclei of helium. This took place until either all of the neutrons had died out and gone forever, or that the particles in the expanding universe were so far apart that they no longer interacted with one another.
One microsecond after the Big Bang all of the neutrinos produced in these reactions were free. They thus became the first fossil relics of the universe. They moved at high speed and their mass, although very small, gives enough gravitational attraction among the hordes that they start clustering together, contributing to the formation of galaxies. About a billion neutrinos are produced for every atom that eventually forms. Neutrinos are thus among the most populous particles in the universe. Although we know that at least one of the varieties of neutrinos has mass, we don’t yet know how big this is. If the mass of a neutrino is greater than a few eV, that is, a billionth of that of a proton, then neutrino masses will dominate the mass density of the material universe. So determining the mass of neutrinos can be a big issue for predicting the long-term future of the universe. Will it expand for ever or eventually collapse under its own weight? We don’t yet know for sure.
The universe continues to expand and cool. The principles of physics that determine its expansion are in some ways similar to those that control the behaviour of a gas in a container. The rate depends on the pressure, which depends on the temperature in the gas and the number of neutrinos inside the gas volume (the density). This in turn depends on the number of neutrino species.
Three minutes after the Big Bang, the material universe consisted primarily of the following: 75% protons; 24% helium nuclei; a small amount of deuterons; traces of other light elements and free electrons.
The abundance of helium and of the light elements depends on the expansion rate of the universe, which in turn depends on the number of neutrino species. The observed amount of helium fits with predictions if there are three varieties of neutrino. The fact that measurements of the Z boson at CERN showed that there are indeed three varieties of light neutrino is a remarkable agreement between measurements in particle physics, which replicates the conditions of the early universe, and what cosmologists had inferred from the above.
The abundance of deuterium depends on the density of ‘ordinary’ matter in the universe (by ordinary we mean made of neutrons and protons as against other exotic things that theorists might dream of but for which there is as yet no direct experimental proof, for example supersymmetry, see Chapter 10). The numbers all fit provided that the density of ordinary matter is much less than the total in the universe. This is part of the dark matter puzzle: there is stuff out there that does not shine but is felt by its gravity tugging the stars and galaxies. It seems that much of this must consist of exotic matter whose identity is yet to be determined.
Some 300,000 years later, the ambient temperature had fallen below 10,000 degrees, that is similar or cooler than the outer regions of our Sun today. At these energies the negatively charged electrons were at last able to be held fast by the electrical attraction to the positively charged atomic nuclei whereby they combined to form neutral atoms. Electromagnetic radiation was set free and the universe because transparent as light could roam unhindered across space.
The universe has expanded and cooled for 10 to 15 billion years so far. The once hot electromagnetic radiation now forms a black body spectrum wit
h an effective temperature of about 3 degrees above absolute zero. The discovery of this by Penzias and Wilson half a century ago is one of the great pieces of support for the Big Bang theory. Today precision measurements of the spectrum by instruments in satellites reveal small fluctuations in the cosmic microwave. These give hints of proto galaxies forming in the early universe.
So we have a good qualitative and even quantitative understanding of how the basic seeds of matter ended up in you and me. But as they emerged along with antimatter in that original Big Bang, a puzzle remains: where did all the antimatter go? At the start of the 21st century, that is one of the questions whose answer is awaited.
Chapter 10
Questions for the 21st century
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Where are we going next? Dark matter in cosmology. Higgs boson – what is it, why do we care, and how might we find it? Precision measurements on exotic heavy particles. Are there more dimensions than those that we presently accept? How might they manifest themselves in experiments? The future of accelerators. Will there be an end to high-energy particle physics?
