Most things around us—including tables, chairs, books, laptops, air, people—are moderately complex when analyzed chemically. The Greeks regarded these various types of matter as mixtures of simple, familiar structures. Combinations of their four basic "elements"—air, earth, fire. and water—were used to describe the nature of everything during much of early recorded history. Later, the Greeks imagined microscopic entities called "atoms," which they regarded as invisible. indivisible, indestructible, and uncreatable substances. This was a step in the right direction.

A step in the wrong direction occurred with the return, during the Dark Ages, to the terrestrially familiar building blocks, including water, dirt, air, oil, sulfur, salt, and so on. Medieval charlatans known as alchemists bamboozled the public by claiming they could change abundant metals such as iron and tin into rare elements such as gold and silver. Not until the 19th century was the concept of atoms resurrected to explain the basic building blocks of matter. Whether this matter was solid, liquid, or gas, these atoms were considered changeless hardly more than 60 years ago.

Figure 3.45 reproduces the familiar Periodic Table of all the known kinds of atoms, or chemical elements, which currently number >100. Arranged in order of increasing complexity from upper left to lower right, the number at the top of each box equals the number of protons or electrons of each element, while that at the bottom equals its total mass. Most of these elements were discovered toward the end of the 18th century and throughout the 19th century.

FIGURE 3.45 FIGURE 3.45 — The Periodic Table of the Elements is a compact way to classify the physical and chemical properties of all known kinds of atoms. For each element, its atomic number is at the top of its chemical symbol and its atomic mass is at the bottom.

Not until well into the 20th century did scientists realize that elements are not changeless. They discovered that nuclei can change from one kind to another. Such nuclear transformations require interactions among elements on scales even smaller than that of atoms; these changes involve nuclear reactions among elementary particles within atoms, rather than chemical reactions among atoms themselves. To some extent, modern nuclear physics has made real the dream of the medieval alchemists. Nuclear laboratories can now change uranium into gold by smashing uranium nuclei with very fast moving subatomic particles. The cost of operating the technology to do so, however, is much greater than the value of the minute amount of gold produced.

The first experimental evidence that elements (really nuclei) can change from one type to another occurred some half-century ago during the basic research that led ultimately to the development of the atomic bomb. Such weapons use uncontrolled fission reactions that break apart heavy nuclei into lighter ones. Similar, though controlled thermonuclear fission reactions in nuclear reactors now power some modern factories, naval vessels, and electrical utilities. More importantly, the opposite process—nuclear fusion, that is, the creation of heavier nuclei from lighter ones like that occurring in ordinary stars or hydrogen bombs—might eventually provide cheap energy needed to power our technological civilization. Indeed, the fusion process considered here is likely to be increasingly important to society on Earth as fossil fuels become depleted in the decades ahead.

Brief Inventory of Matter First, some terminology to clearly distinguish atoms from ions from isotopes from nuclei from elementary particles. Recall that every positively charged atomic nucleus is extremely small compared to the size of the cloud of negatively charged electrons orbiting about it. While in a hydrogen atom, a single electron orbits a single proton, heavier elements quickly become more complex with neutrons coexisting in the nucleus alongside the usual protons, around which many electrons move in discrete orbitals. When an atom has a net charge because the numbers of electrons and protons are unequal, we call this charged atom an ion.

Atoms of different kinds—that is, elements—are distinguished by their proton number or their electron number. In all, we know of some 107 different elements, from the simplest, hydrogen (1 proton), to the most complex, copernicium (112 protons). (Heavier elements 114 and 115 have been confirmed but are as yet unnamed, while claims for elements having 113, 115, and 118 remain unconfirmed.) Each element usually has a fixed number of protons and neutrons within its nucleus. These are the normal atoms comprising baryonic matter noted earlier in the PARTICLE EPOCH.

Sometimes, however, the number of neutrons in a given element differs slightly from its standard value. These are rarer atoms termed "isotopes". Having a few more or less neutrons than usual, isotopes play important roles in the creation of the heavy elements. Many isotopes are radioactively unstable, hence decay into more stable isotopes or ordinary atoms after specific amounts of time.

Table 3-3 lists a brief inventory of the currently known atomic building blocks of nature. The 81 stable elements comprise the bulk of the list. All of them are found in Earth's air, land, and sea.

In addition, 10 radioactive elements naturally reside on our planet. These elements are termed "natural" because they’re created naturally (actually in stars). Like all radioactive elements, they have certain half-lives, after which half of the total amount of that element decays into something else. Even though their half-lives are often very long (typically millions to billions of years), the steady decay or disappearance of these natural radioactive elements explains their scarcity on Earth, in meteorites, and in lunar samples. They’re not observed in stars because their abundances are too low to produce detectable spectral lines.

In turn, ~14 more radioactive elements can be artificially produced under special conditions in nuclear laboratories on Earth. The debris collected after nuclear weapons tests also show evidence for these artificial radioactive elements. Unlike the naturally occurring radioactive elements, these artificial (human-made) ones often decay into other elements quickly (typically much less than a million years); the newest one, copernicium named after the 16th century astronomer Copernicus, lasts for less than a second. Consequently, we shouldn’t be surprised that they’re extremely rare or virtually nonexistent in Nature, except presumably in supernovae where they’re occasionally created (see below).

Rounding out this list of elements, another one—promethium—forms only as a by-product of nuclear laboratory experiments. And still another unstable element—technetium—is found only in stars. Neither of these occurs naturally on Earth.

Matter even at the atomic level is not so simple, however. There are many more isotopes than elements. The element carbon, for example, is known to have 10 isotopes, 8 of them radioactive. In all, scientists know of nearly 300 stable isotopes, ~70 naturally occurring radioactive isotopes, and >2700 artificially produced radioactive isotopes. The total list of isotopes exceeds 3000 and grows longer each year as nuclear researchers discover more rare and exotic isotopes of the known elements.

Are these ~3000 elements and isotopes the absolute basic building blocks of matter? The answer is "no," of course. We know that atoms’ more basic entities—the electrons, protons, and neutrons—are, in turn, made of smaller, even more basic particles—the quarks noted in the PARTICLE EPOCH. Physicists are currently probing the composition of protons and neutrons by smashing them together in the world’s most powerful accelerators. Although the resultant debris should tell us something about the internal makeup of these subatomic particles, the results of these experiments aren’t yet clear. Fortunately, to understand the origin of the elements, we need not worry about the detailed structure of matter below the level of the proton, neutron, and electron.

Abundance of Matter Today’s understanding of heavy-element creation is aided by experiments in nuclear laboratories that have proven that larger elements can be built from smaller ones. This is usually accomplished by fusing together two or more relatively light nuclei in a violent collision. Among the collision's debris of energy, elementary particles and light nuclei, are some heavier nuclei as well. Hence, we theorize that all the heavy elements were built up from the lighter elements. In this scheme, the ultimate source of the heavy elements is the lightest and simplest one—hydrogen.

To test this idea, consider not only the list of different kinds of elemental atoms and isotopes (Table 3-3 ), but also the observed abundances of these elements. The cosmic abundance scale is shown in Figure 3.46, derived largely from the spectroscopy of many stars including the Sun. Table 3-4 is another way of summarizing the essence of this key figure. (Various isotopes of all the elements are also included in both Figure 3.46 and Table 3-4 .) Any theory proposed for the creation of the elements must match these observed abundances, the most obvious feature being that the heavy elements are much less abundant than most light elements.

FIGURE 3.46 FIGURE 3.46 — A summary of the cosmic abundances of the elements and their isotopes. These are general abundances of the most common material entities in the Universe, namely stars and interstellar matter; they do not match the abundances of rocky planets and their moons, which are minor players in the cosmic scheme of things.

Approximately 10 hydrogen atoms exist for every 1 helium atom. Together, these two lightest atoms comprise ~99% of all matter in the observable Universe. Next comes a group of slightly heavier elements, namely those having from 7 to 11 particles in their nuclei: lithium, beryllium, and boron. Even summed together, this so-called lithium group has an extraordinarily small abundance. Slightly heavier elements, such as carbon, nitrogen, and oxygen, are more abundant than this lithium group, though still much less abundant than the lightest hydrogen and helium elements. The same can be said for the silicon and iron groups listed in Table 3-4 . The middle-weight and heaviest-weight groups—those with 63 or more nuclear particles—have progressively lesser abundances. To the annoyance of most chemists, astronomers refer to all those heavier than hydrogen and helium as “heavy elements,” or simply “heavies.”

These relative abundances show an interesting peculiarity. Despite the overwhelmingly large amount of hydrogen and helium throughout the Universe, not much of either of these elements exists on Earth. That's because the Terrestrial Planets don’t typify the average matter in the Universe. For the most part, the various abundances of rocky planets don’t mimic those of stars or galaxies. Earth's small gravity and large solar heating combine to make our planet rather inhospitable for hydrogen and helium (<1% on average on Earth), the result being that these light gases don’t stick around. Instead, most of the elements surrounding us are of intermediate mass—much silicon in rocks, nitrogen in air, and carbon in people (the most abundant Earth element by mass is actually iron, most of it at the core of our planet). Planets resemble “islands” of abnormal abundances in a giant sea of normal cosmic abundances. This is especially true for the rocky Terrestrial Planets, less so for the gassy Jovian Planets. The bigger and more distant planets of our Solar System have stronger gravity and less solar heating than the smaller, interior planets, and thus have elemental abundances more like those listed in Table 3-4 .

Primordial Nucleosynthesis The observed cosmic abundances can be reconciled in two ways. First, we can theorize that all the elements were created shortly after the initial event that started the Universe, in which case the entire process is called primordial nucleosynthesis . Or, we can argue that the heavies are produced later within stars, in which case the process is called stellar nucleosynthesis . Let's consider each of these possibilities in turn.

Imagine the physical conditions just after the origin of the Universe. As described in the PARTICLE EPOCH, the temperature of matter and the density of radiation were very high. Neutrons were abundant, as were many other elementary particles, all racing out from the violent expansion. However, a free neutron existing alone in space (neither as part of an atom nor combined in some other structure such as a neutron star) cannot remain that way for long. In ~11 minutes, it decays into an electron and a proton, a measured fact:

neutron --> electron + proton + neutrino.

Excepting the neutrino particle, note that the products comprise hydrogen. Although the early Universe was too hot for electrons and protons to couple together and actually form atoms, it’s interesting to realize that hydrogen nuclei first came into being in this way. Given the seething conditions of those first few minutes, this proton would have nonetheless combined with another free neutron (before it decayed) in order to produce a deuteron particle:

proton + neutron --> deuteron.

A deuteron is nothing more than the nucleus of deuterium, a heavier, isotopic form of hydrogen. Deuterium is an isotope because it contains one extra neutron than does normal hydrogen, but no additional protons. This deuteron particle can then in turn combine with another free neutron to form a triton particle:

deuteron + neutron --> triton,

which is the nucleus of the heaviest form of hydrogen, another isotope called tritium. Like free neutrons, triton particles are unstable and hence decay into a special form of helium:

triton --> 3helium,

which, in turn yet again, can then interact with another free neutron, thus becoming a more stable form of helium:

3helium + neutron --> 4helium.

In addition to repeated neutron capture of this sort, protons could also have been engaged to yield both forms of helium. Such "captures" of neutrons and protons would have almost certainly occurred in the first few minutes of the Universe. These reactions are thus feasible ways of synthesizing hydrogen and helium nuclei. Further capture of electrons by these nuclei later in time, once the Universe calmed down a bit, would have produced the lightest elements.

All the above nuclear events did almost surely occur in the early PARTICLE EPOCH. But several problems arise when we postulate these kinds of nuclear events to understand the creation of the heavy elements. One issue is that the decay of the triton particle into the 3helium nucleus takes ~12 years. That's a short time in the cosmic scheme of things, but in the earliest parts of the Universe the physical conditions must have changed very rapidly. The Universe expanded greatly during the first decade after the bang, severely diluting the density of neutrons. Accordingly, the last reaction noted above, namely neutron capture by isotopic 3helium to form ordinary 4helium, is improbable. No doubt this nuclear event did occur early on, but only in a few places and rather infrequently at that.

Another problem with continuing the scenario above to produce heavies is that, among all those elements found in Nature, elements of mass 5 and 8 are missing. We know of no stable atoms or isotopes having those mass values, and these represent clear and obvious gaps at those positions in the Periodic Table that must be explained. If the neutron capture scheme continued on indefinitely, it would be difficult to create elements much more complex than helium if elements of mass 5 and 8 were skipped. More basically, if Nature built up the heavies by capturing neutrons one at a time, why would it not have made those non-existent, lightweight masses?

A third objection impairs primordial nucleosynthesis, and this one is not merely theoretical. If the heavies were created by repeated neutron capture in the early Universe, all stars should then have identical elemental abundances. Reality shows otherwise, however, for we observe different abundances in different objects. On average, all stars have roughly the same cosmic abundances as those listed in Table 3-4 , but some of them do have small, yet distinct elemental differences. This is especially true when comparing the young stars in galactic clusters that are rich in heavy elements with the old stars in globular clusters that are mostly lacking in heavy elements.

A fourth problem concerns the temperature and density of the early Universe. When heavy elements are made via any kind of nuclear fusion event, the electromagnetic charge of their nuclei inevitably increases. Progressively greater force is then needed to combine the heavies into even heavier nuclei. Higher temperatures and greater densities would normally be able to do it, but the early Universe was doing just the opposite—rapidly cooling and thinning, as depicted in Table 1-1. Thus, the process of primordial nucleosynthesis, which requires increasing temperatures and densities, is inconsistent with the trend of an expanding Universe where these quantities were decreasing.

So many problems plague the heavy-element scenario of primordial nucleosynthesis that it doesn’t seem viable. This process would have had time to form some helium (and perhaps traces of the lightest-weight lithium group), but not enough time to form any of the heavier elements. Events at the start of the Universe happened too rapidly.

Stellar Nucleosynthesis Astronomers now realize that better places for the creation of heavy elements are the stars themselves. In the interiors of stars, both temperature and density are high, and get progressively higher as stars evolve. In short, stars are pockets of increasing temperature and density within a Universe of decreasing temperature and density.

The production of heavies in stars—stellar nucleosynthesis—seems more reasonable not only because primordial nucleosynthesis is riddled with problems, but also because of much favorable theoretical insight and observational evidence. For example, the ages of stars and their evolutionary paths suggest that heavies are now being created deep inside stars—and spectroscopy confirms this idea. Also, differences in heavy-element abundances between the older globular-cluster stars and the younger galactic-cluster stars clearly imply that the heavies are slowly produced over the course of time. Furthermore, stars are the only places known where temperatures and densities remain high over long durations, allowing the heavies to be processed steadily well after the start of the Universe.

Stellar nucleosynthesis relies on the usual stellar burning cycle of contracting, heating, expanding, cooling, contracting, and so on, thereby making the heavies during each successive period of stellar contraction. It all begins with the so-called proton-proton cycle. Provided the temperature and density are suitable—at least 107 K and 100 g/cm3—the following nuclear reaction proceeds automatically:

1hydrogen + 1hydrogen --> deuteron + positron + neutrino + energy.

(Remember, these are nuclear reactions where the symbols refer to nuclei only, not chemical reactions involving whole atoms. 1Hydrogen, for example, represents a proton; deuteron is a blend of proton and neutron. Once nuclei are created, they can later capture appropriate numbers of electrons in order to form different kinds of atoms.) The positron particle produced here immediately interacts with a nearby free electron, thereby producing high-energy radiation via matter-antimatter annihilation. The neutrino particle rapidly escapes the scene, carrying with it some energy, but playing no direct role in stellar nucleosynthesis. Only the deuteron sticks around to participate in further reactions.

In stars, ordinary helium is made by means of a two-step process, its validity directly confirmed in nuclear experiments conducted in laboratories around the world during the past few decades. First, an isotope of helium is formed in the nuclear reaction, deuteron + 1hydrogen --> 3helium + energy, after which further interaction of these isotopic nuclei produce the normal form of helium:

3helium + 3helium --> 4helium + 1hydrogen + 1hydrogen + energy.

The use of 6 hydrogen nuclei (protons) and the return of 2 of them along with a 4helium nucleus means that 4 hydrogen nuclei are needed to create a single 4helium nucleus. This is hardly surprising given that hydrogen has a 1-particle nucleus (a proton), while 4helium has a 4-particle nucleus (2 protons and 2 neutrons).

The combined mass of the reactants always exceeds that of the products, ensuring that energy is produced during these nuclear-fusion reactions—all in accord, yet once more, with Einstein’s famous equation, E = mc2. Some of this energy is the heat that counteracts gravity to prevent stars from collapsing. The rest is radiation that escapes the star as starlight.

As a star steadily increases the amount of helium within its central regions, its core gradually cools for lack of nuclear burning among the helium nuclei. The inward pull of gravity slowly overwhelms the diminished heat (gas pressure) pushing outward. The core then contracts a little, increasing the density of nuclei to a value of ~100,000 g/cm3. More frequent collisions among the many nuclei then build up the heat again, until the temperature reaches ~108 K. Such hot and dense physical conditions automatically ignite the fusion of nuclei according to the following nuclear reaction:

4helium + 4helium --> 8beryllium + energy.

The elementary particles in a 8beryllium nucleus don’t attract each other very strongly, however. As an unstable isotope, it decays almost immediately (in a trillionth of a second) into a more stable nucleus. This extraordinarily short lifetime of 8beryllium accounts for the fact that, as noted above, there’s no element number 8 in the Periodic Table (Figure 3.45); beryllium doesn’t naturally persist in Nature. Nonetheless, the 8beryllium isotope has been detected for fleeting moments in laboratory accelerators, so physicists know a little bit about it.

Deep within the cores of aged stars, the density is so great that some 8beryllium nuclei can collide with other nuclei before changing into something else. After all, stars are so large that vast quantities of helium are fusing all the time. At any given moment, small amounts of 8beryllium are likely just produced but not yet decayed. If a helium nucleus happens to collide with an intact 8beryllium nucleus before it has had a chance to decay, a stable nucleus of carbon is created:

8beryllium + 4helium --> 12carbon + energy.

Note that the net result thus far is for 3 4helium nuclei to synthesize 1 12carbon nucleus. This again makes sense, for 3 nuclei containing 2 protons and 2 neutrons apiece naturally ought to form 1 nucleus containing 6 protons and 6 neutrons.

This type of helium-capture process continues in order to construct heavier elements. For example, provided that the temperature is at least 6x108 K, any 12carbon nucleus colliding violently with another 4helium nucleus produces a stable oxygen nucleus:

12carbon + 4helium --> 16oxygen + energy.

With each expansion-contraction cycle of a star, the temperature and density rise at the core, enabling 4helium nuclei to be captured by progressively heavier nuclei. If the temperature didn’t grow larger, that is, if the velocities of the nuclei did not increase, the particle collisions wouldn’t be violent enough to overwhelm the progressively stronger electromagnetic repulsions of the heavier nuclei. Fortunately, with time, gravity increasingly contracts stellar cores, causing higher and higher temperatures. In this way, heavier elements such as 20neon, 24magnesium, and 28silicon are eventually produced, along with more energy at all stages, as stars move far from the main sequence on the HR diagram. Nucleosynthesis becomes relatively straightforward especially for those nuclei divisible by 4. Because such "helium capture" reactions are so frequent, elements numbered 12, 16, 20, 24, and 28 stand out as abundant "peaks" in the Figure 3.46 chart of cosmic abundances.

Helium capture is not the only type of nuclear reaction occurring in stars. As more nuclei of different kinds accumulate in aged stars, a great variety of nuclear reactions become possible. Some nuclei capture neutrons, protons, and deuterons, in addition to helium nuclei. The result is a family of many nuclei having masses intermediate to those mentioned above. In fact, laboratory studies confirm that nuclei such as 19fluorine, 23sodium, 31phosphorus, and a wide variety of others are created by the steady capture of small increments of mass and charge. Their abundances, however, aren’t as great as those nuclei produced directly by helium capture. That's why many of these elements (not divisible by 4) reside in the "troughs" of the Figure 3.46 chart of cosmic abundances.

Some Complications With the appearance of 28silicon in the core of a star, another nuclear process begins to complicate the rather simple and repetitive 4helium capture scheme. A competitive struggle ensues between the continued capture of helium nuclei to produce even heavier nuclei and the tendency of the heavier nuclei to break down into simpler ones. The cause of this breakdown, called “photodisintegration,” is simply heat, indeed lots of it. By this point in a star’s evolution, it’s core temperature has reached an unimaginably large value of ~3 billion kelvins, much, much higher than anything we can spark on Earth. The density is also enormous at this time, having millions of g/cm3. Even sturdy nuclei have difficulty remaining intact within such an absolutely raging inferno.

Consider an example of typically constructive and destructive events thought to occur in highly evolved stars. Heat is so intense that some of the 28silicon nuclei break apart into 7 4helium nuclei. Another nearby 28silicon nucleus, not yet itself photodisintegrated, then suddenly captures one or several of these 4helium nuclei. This type of process leads to even heavier nuclei including 32sulfur, 36chlorine, 40argon, 40calcium, 48titanium, and 52chromium—again, all at intervals of 4. In the rare case where all 7 helium nuclei from a photodisintegrated silicon nucleus are suddenly captured by a nearby undestroyed 28silicon nucleus, a much more massive nucleus of 56nickel is instantaneously created:

28silicon + 4helium + 4helium + 4helium + 4helium + 4helium + 4helium + 4helium --> 56nickel + energy.

This two step process—photodisintegration followed by direct capture of disintegrated 4helium nuclei—is called the alpha process, a name again derived from the shorthand name for 4helium nuclei.

A further complication enters the picture here. The most stable nickel nucleus is 60nickel, not 56nickel as produced in this reaction. The 56nickel isotope then quickly decays first into a 56cobalt isotope and then into a normal 56iron nucleus. As a rule of thumb in Nature, any unstable nucleus will continue to decay until stability is achieved. We know this much from experiments in nuclear physics laboratories. And no known nuclei are more stable than 56iron.

Iron's 26 protons and 30 neutrons are bound together more strongly than the particles in any other nucleus; as noted earlier, iron has the highest nuclear binding energy. Any nucleus with greater or fewer protons or neutrons has less binding energy and thus cannot be quite as stable as 56iron. This enhanced stability of iron nuclei, as well as of several others having similar complexity, explains how some of the heavier nuclei in the iron group have become more abundant than several lighter nuclei. That this is indeed the case can be seen by noting the cosmic abundance of iron and its neighbors in Figure 3.46.

The unusual stability of the 56iron nucleus also endows it with another important property—it’s unable to fuse with other nuclei to produce more energy. In fact, fusion events involving iron consume energy. This is a fundamental change in the overall trend of stellar nucleosynthesis. Nuclear fusion in a highly evolved massive star (much more massive than the Sun) inevitably produces an iron core that not only fails to produce energy, but also begins to rob the star of some of its previously gained energy. This is the sense, as noted earlier, that iron nuclei act as fire extinguishers, quenching nuclear fires in the stellar cores—at which point the most massive stars start to panic, unstablize, and collapse.

Stars reaching the stage at which iron is synthesized have long left the main sequence on the HR diagram. Nuclear burning is under way at all interior layers of such stars. Figure 3.47 is a cutaway painting of such a fully involved star, now truly resembling a multi-layered onion. The peripheral stellar layers produce helium via the basic proton-proton cycle, the intermediate layers synthesize a variety of heavier-than-helium nuclei mainly via helium-capture events, and the innermost core layers house the sites of the alpha-process that creates heavy nuclei up to and including iron.

FIGURE 3.47 FIGURE 3.47 — Artist's conception of the interior of a highly evolved massive star, showing as a cutaway diagram the many internal layers of nuclear burning. (Lola Chaisson)

Heaviest Elements None of the above mechanisms create nuclei heavier than iron. Electromagnetic repulsion between helium and heavy nuclei, all of them positively charged, is too great for fusion to occur even in the inferno-like cores of aged stars. But if the heaviest nuclei at the bottom of the Periodic Table of Elements are to be synthesized—and clearly they must be somehow made because we know they exist, in the coins in our pockets, the jewels on our fingers, among other precious metals valued by society—then other nuclear events must be at work. Fortunately, another such process works just fine, this one involving the capture of uncharged neutrons.

Deep in the hearts of stars that are massive, highly evolved, and collapsing, the conditions are ripe for the elementary particles to play a role once again. Their numbers swell fantastically as nuclei succumb to the onslaught of heat and pressure during catastrophic collapse. Neutrons, in particular, can easily interact with other heavy nuclei that haven’t yet broken down; without any charge, neutrons have no electromagnetic repulsive barrier to overcome. For example, an iron nucleus can successfully capture a single neutron to form a relatively stable isotope, 57iron:

56iron + neutron --> 57iron.

This is then quickly followed by another neutron capture, 57iron + neutron --> 58iron, producing a relatively stable 58iron isotope, which can capture yet another neutron to produce an even heavier isotope:

58iron + neutron --> 59iron.

59Iron, however, is known from laboratory experiments to be radioactively unstable. It decays in about a month into 59cobalt, after which this neutron-capture process resumes. In this way, further neutron captures can progressively form much heavier nuclei, including those of silver and tin in some monetary coins, or gold in rings and watches. Such nuclear reactions continue all the way up to 209bismuth, which is the heaviest product possible in this scheme of successive neutron captures—the heaviest nonradioactive nucleus known:

59cobalt + neutron --> 60cobalt................

60cobalt (decay) --> 60nickel................












We might be surprised to learn that this neutron-capture method isn’t very fast by nuclear physics standards. Nuclei don’t rapidly cascade through the entire series of reactions that change them from iron to bismuth, even though the tremendously hot and dense cores of massive stars would suggest a feverish pace. Instead, each capture of a neutron by a nucleus typically takes about a year. Consequently, researchers refer to this "slow" neutron-capture mechanism as the “s-process .”

As noted earlier in this STELLAR EPOCH, the long chain of stellar nucleosynthesis is bolstered by two pieces of observational evidence: Predictions of the relative amounts of each kind of nucleus produced matches remarkably well the observed cosmic abundances (Table 3-4 )—all the way up to and including iron. Additionally, observations in red giants of one kind of nucleus—99technetium—provides direct evidence that heavy-element formation really does occur in stars today.

The s-process goes beyond iron, allowing us to understand the workings of nucleosynthesis all the way up to 209bismuth. But we've yet to reach some of the heaviest nuclei such as those of 232thorium, 238uranium, and 242plutonium. Since the s-process terminates at bismuth, there must be yet another nuclear mechanism that produces the heaviest nuclei of all. There is indeed such another process known—the “r-process,” where "r" stands for "rapid." Another neutron-capture scheme, the r-process operates quickly because it occurs literally during the moment of explosion as a massive star undergoes a supernova outburst.

Supernovae rebound like giant coiled springs after the nuclei themselves halt the stellar collapse, spewing forth a multitude of heavy elements created within. At the instant of explosion and for ~15 minutes thereafter, the density of neutrons dramatically increases as some heavy nuclei break apart during the explosion itself. Jammed into many of the light- and middle-weight nuclei, the neutrons help create the heaviest of the known nuclei. Interestingly enough, then, the heaviest of the heavies are actually produced after stars have died.

Because the time available for synthesizing the heaviest of all nuclei is so brief, not surprisingly they never become as abundant as nuclei up to and including iron. The cosmic abundances (Table 3-4 ) of elements heavier than iron are a billion times less abundant than hydrogen and helium.

Again, astronomers have some evidence that this r-process really does occur in Nature—at least indirect evidence. Figure 3.48 is a typical light curve of a supernova, displaying its dramatic rise in luminosity signifying the moment of explosion, followed by a characteristically slower decay. Depending on the initial mass of the exploded star, the luminosity might take from a few months to many years to decrease to its original value. But the shapes of the decay curves are pretty much the same for all exploded stars.

FIGURE 3.48 FIGURE 3.48 — The light curve of an actual supernova, showing not only the dramatic increase and slow decrease in luminosity, but also the characteristic change in the rate of decay (arrow) ~2 months after the explosion. (This supernova, labeled IC4182, was observed in a faraway galaxy in 1938.) (Harvard College Observatory)

Detailed analyses show these decay curves to have two distinct features. As can be seen from Figure 3.48, the post-peak luminosity first decreases rapidly, after which it continues decreasing but at a slower rate. This change in the luminosity decay invariably occurs ~2 months after the explosion. Regardless of the supernova or the intensity of the outburst, the break in the decay curve always seems to happen ~50 days after the explosion. The peculiar twofold nature of this curve is not well understood, but many researchers suspect the answer lies in the rubble of hydrogen (fusion) bomb tests. Indeed, the rise, decay, and overall shape of the luminosity curve of a hydrogen bomb blast mimics quite closely that of most supernovae.

Figure 3.49 plots the decay of radioactivity measured for a group of elements created during human-made thermonuclear explosions here on Earth. Several curves are shown, each one representing a different heavy element with an unstable radioactive nucleus. Note, for example, the decay curve of the highly radioactive nuclei 56nickel and 56cobalt. These unstable isotopes, found in abundance among the debris after nuclear bomb tests, have half-lives of 6 days and 78 days, respectively; together, their decay (and thus their release of energy) mimic the light curves observed toward supernovae. In fact, a gamma-ray spectral feature attributed to radioactive cobalt was tentatively identified in a supernova that occurred in a distant galaxy a few decades ago. Other, even heavier, radioactive isotopes are also produced in nuclear explosions, although most have longer decay times measured in years.

Thus some astronomers reason that the observed luminosity curves of supernovae can be explained by accumulating the radioactive decays of many unstable nuclei. To be sure, the similarity in the luminosity changes of hydrogen bomb blasts and of supernova explosions are considered more than coincidence. It's taken as good indirect evidence that the r-process really does work immediately folllowing the death of the most massive stars.

FIGURE 3.49 FIGURE 3.49 — The characteristic radioactive decay curves for some very heavy nuclei found among the debris of hydrogen bomb tests on Earth. (US Defense Dept.)

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