Nuclear Astrophysics – The Synthesis of the Chemical Elements
Professor Zsolt Podolyak, University of Surrey
Professor Podolyak was born in Romania and obtained his PhD in Hungary. He has worked in Italy and is now based at the University of Surrey, where among other topics teaches Nuclear Astrophysics. He is a Nuclear Physicist, with research based on measurements at international facilities in Europe, USA and Japan.
The creation of chemical elements in nature, from the lightest element (hydrogen) to the heaviest (Uranium) was discussed in the talk. Different processes, some slow (lasting for many billions of years), some very fast (such as supernova explosions) were introduced. The relevance of nuclear physics and ways to measure the relevant quantities were covered.
The age of the Earth
In around 1890 Lord Kelvin estimated that the Earth was about 20 to 40 million years old. He was a devout Christian and saw his Christian faith as supporting and informing his scientific work. One of the clearest instances of this interaction was in his estimate of the age of the Earth.
Given his youthful work on the figure of the Earth and his interest in heat conduction, it is no surprise that he chose to investigate the Earth’s cooling and to make historical inferences of the Earth’s age from his calculations. Thomson was a creationist in a broad sense, but he was not a ‘flood geologist’. He contended that the laws of thermodynamics operated from the birth of the universe and envisaged a dynamic process that saw the organisation and evolution of the solar system and other structures, followed by a gradual “heat death”. He developed the view that the Earth had once been too hot to support life and contrasted this view with that of uniformitarianism, that conditions had remained constant since the indefinite past. He contended that “This earth, certainly a moderate number of millions of years ago, was a red-hot globe.
He was not a supporter of Charles Darwin and favoured a version of theistic evolution sped up by divine guidance. His calculations showed that the Sun could not have possibly existed long enough to allow the slow incremental development by evolution – unless some energy source beyond what he or any other Victorian era person knew of was found..
His initial 1864 estimate of the Earth’s age was from 20 to 400 million years old using the idea of thermal gradients. These wide limits were due to his uncertainty about the melting temperature of rock and his not knowing about the highly viscous fluid mantle, to which he equated the earth’s interior temperature. Over the years he refined his arguments and reduced the upper bound by a factor of ten, and in 1897 he ultimately settled on an estimate that the Earth was 20–40 million years old.
William Thomson, 1st Baron Kelvin OM GCVO PC PRS PRSE (26 June 1824 – 17 December 1907) was a British mathematical physicist and engineer who was born in Belfast in 1824.
Charles Robert Darwin, FRS (12 February 1809 – 19 April 1882) was an English naturalist and geologist, best known for his contributions to evolutionary theory.
The discovery in 1903 that radioactive decay releases heat led to Kelvin’s estimate being challenged, and Ernest Rutherford famously made the argument in a lecture attended by Kelvin that this provided the unknown energy source that Kelvin had suggested, but the estimate was not overturned until the development in 1907 of radiometric dating of rocks.
Ernest Rutherford, 1st Baron Rutherford of Nelson, OM, FRS (30 August 1871 – 19 October 1937) was a New Zealand-born British physicist who became known as the father of nuclear physics.
In his 1905 lecture Rutherford had proposed that the Earth was billions of years old and he suggested to Bertram Boltwood that he investigate the decay of uranium to lead.
Bertram Borden Boltwood (July 27, 1870 Amherst, Massachusetts – August 15, 1927, Hancock Point, Maine) was an American pioneer of radiochemistry.
He graduated from Yale University, and taught there 1897-1900. He established that lead was the final decay product of uranium, noted that the lead-uranium ratio was greater in older rocks and, acting on a suggestion by Ernest Rutherford, was the first to measure the age of rocks by the decay of uranium to lead, in 1907. He got results of ages of 400 to 2200 million years, the first successful use of radioactive decay by Pb/U chemical dating (isotopes not discovered yet). More recently, older mineral deposits have been dated to about 4.4 billion years old, close to the best estimate of the age of earth.
The Milky Way
Nucleosynthesis in cosmic sources can be observed through a variety of, mostly indirect, measurements; examples are stellar photospheric absorption lines, or mass spectrometry of meteoritic inclusions. Gamma-rays from radioactive by-products of nucleosynthesis ejecta provide a rather direct measurement, in comparison, as their decay gamma-ray measurements with satellite-borne telescopes provide direct isotopic constraints to the physics of nuclear burning regions inside these sources. Yet, the technique of gamma-ray telescopes is complex, and less precise than the alternatives for cosmic abundance measurements, mainly from two reasons: Spatial resolutions of ~degrees and signal-to-background ratios of ~% restrict contributions from gamma-ray astronomy to nearby sources in the Galaxy. There are, however, advantages to gamma-ray astronomical data: They provide isotopic information, are unaffected by physical conditions in/around the source such as temperature or density, and gamma-rays are nearly un-attenuated along the line-of-sight due to their penetrating nature (attenuation length ~few g/cm^2). Longer-lived isotopes such as 26Al (t~1.04 Myrs) will reflect the kinematic properties in an interstellar medium around massive star sources, which is otherwise hard to study due to its hot and diluted nature.
The image above is a gamma picture of the galaxy obtained from the decay of 26Al (an isotope of aluminium not found on Earth). The half-life of the decay is 0.74Myear and the energy of the gamma rays is 1.8MeV
A scheme of the science goals: Radioactive decay as measured by INTEGRAL from 26Al decay, and the signature of Galactic rotation. The background image of the Milky Way is overlaid with the COMPTEL map of 26Al emission. Reference: MPE, Oct 2005
The gamma emission at 1809 keV was the first observed gamma emission from the galactic centre. The observation was made by the HEAO-3 satellite in 1984.
The isotope is mainly produced in supernovae ejecting many radioactive nuclides in the interstellar medium. The isotope is believed to provide enough heat to small planetary bodies so as to differentiate their interiors, such as has been the case in the early history of the asteroids 1 Ceres and 4 Vesta. This isotope also features in hypotheses regarding the equatorial bulge of Saturn’s moon Iapetus.
The gamma-rays from radioactive 26Al provide a snapshot view of continuing nucleosynthesis in the Galaxy. The 26Al is produced by nucleosynthesis is created in various processes in stellar evolution.
Nucleosynthesis is the process that creates new atomic nuclei from pre-existing nucleons, primarily protons and neutrons. The first nuclei were formed about three minutes after the Big Bang, through the process called Big Bang nucleosynthesis. It was then that hydrogen and helium formed to become the content of the first stars, and this primeval process is responsible for the present hydrogen/helium ratio of the cosmos.
With the formation of stars, heavier nuclei were created from hydrogen and helium by stellar nucleosynthesis, a process that continues today. Some of these elements, particularly those lighter than iron, continue to be delivered to the interstellar medium when low mass stars eject their outer envelope before they collapse to form white dwarfs. The remains of their ejected mass form the planetary nebulae observable throughout our galaxy.
Supernova nucleosynthesis within exploding stars by fusing carbon and oxygen is responsible for the abundances of elements between magnesium (atomic number 12) and nickel (atomic number 28). Supernova nucleosynthesis is also thought to be responsible for the creation of rarer elements heavier than iron and nickel, in the last few seconds of a type II supernova event. The synthesis of these heavier elements absorbs energy (endothermic) as they are created, from the energy produced during the supernova explosion. Some of those elements are created from the absorption of multiple neutrons (the R process) in the period of a few seconds during the explosion. The elements formed in supernovas include the heaviest elements known, such as the long-lived elements uranium and thorium.
Cosmic ray spallation, caused when cosmic rays impact the interstellar medium and fragment larger atomic species, is a significant source of the lighter nuclei, particularly 3He, 9Be and 10, 11B, that are not created by stellar nucleosynthesis.
In addition to the fusion processes responsible for the growing abundances of elements in the universe, a few minor natural processes continue to produce very small numbers of new nuclides on Earth. These nuclides contribute little to their abundances, but may account for the presence of specific new nuclei. These nuclides are produced via radiogenesis (decay) of long-lived, heavy, primordial radionuclides such as uranium and thorium. Cosmic ray bombardment of elements on Earth also contributes to the presence of rare, short-lived atomic species called cosmogenic nuclides.
Look at isotopes for evidence of nucleosynthesis.
What are things made of?
Air consists of 78% nitrogen, 21% Oxygen and 1% of trace elements such as argon. A pebble on the beach consists of 47% silicon, 53% oxygen (by weight) plus traces of elements such as iron and manganese.
What is the Solar System made of?
The Oddo–Harkins rule holds that elements with an even atomic number (such as carbon) are more common than elements with an odd atomic number (such as nitrogen). This effect on the abundance of the chemical elements was first reported by Giuseppe Oddo in 1914 and William Draper Harkins in 1917.
The graph above shows the estimated abundances of the chemical elements in the Solar system. Hydrogen and helium are most common, from the Big Bang. The next three elements (Li, Be, B) are rare because they are poorly synthesized in the Big Bang and also in stars. The two general trends in the remaining stellar-produced elements are: (1) an alternation of abundance in elements as they have even or odd atomic numbers (the Oddo–Harkins rule), and (2) a general decrease in abundance, as elements become heavier. Iron is especially common because it represents the minimum energy nuclide that can be made by fusion of helium in supernovae. Note: element 43 (Tc) Technetium is missing from the graph due to its extremely low abundance. This appears to, but doesn’t actually, throw off the rule if one is just looking at the ups and downs on the graph.
The Solar System is 74% hydrogen, 24% helium and 2% other elements including carbon and oxygen.
Elemental Abundance in the Universe
Hydrogen (75%) and helium (23%) are the most abundant elements in the universe. Other abundant elements like carbon and oxygen were finally produced in later generations of stars. Stars make elements throughout their life and contribute them to the interstellar space at death. Notice the dominance of hydrogen in the solar system in the following chart. Most of the elements that comprise our bodies are “impurities” in the solar system when compared to hydrogen and helium.
Elements and isotopes
An element is a substance made from only one type of atom and contains a unique number of protons.
Isotopes are variants of a particular chemical element which differ in neutron number, although all isotopes of a given element have the same number of protons in each atom.
The number of neutrons is designated N and the number of protons is designated Z.
The chemical properties of an element (and their isotopes) are dictated by the number of protons (Z).
The physical properties of an isotope are dictated by the sum of the protons and neutrons in the nucleus (A = N + Z).
The nuclear properties and synthesis of elements are dictated by the nuclei (N + Z)
The Periodic Table
The periodic table is a tabular arrangement of the chemical elements, ordered by their atomic number (number of protons in the nucleus), electron configurations, and recurring chemical properties.
In nuclear physics, a magic number is a number of nucleons (either protons or neutrons) such that they are arranged into complete shells within the atomic nucleus. The seven most widely recognized magic numbers as of 2007 are 2, 8, 20, 28, 50, 82, and 126. Atomic nuclei consisting of such a magic number of nucleons have a higher average binding energy per nucleon than one would expect based upon predictions such as the semi-empirical mass formula and are hence more stable against nuclear decay.
The unusual stability of isotopes having magic numbers means that transuranium elements can be created with extremely large nuclei and yet not be subject to the extremely rapid radioactive decay normally associated with high atomic numbers. Large isotopes with magic numbers of nucleons are said to exist in an island of stability.
Chart of nuclei
The black bits in the chart are found in nature. The yellow bits are synthesised. The unknowns are involved in nucleosynthesis.
Far from the line of stability you should be able to create nuclei where it is energetically favourable to just eject the proton or neutron completely. This is because as you move away from the line of stability you are raising the potential energy of the nucleons without raising the binding force as you would if you add more nucleons. The lines that define this region in the N vs. Z planes go from the region of stability where you see more common isotopes and decays and extend to lines known as drip-lines. The drip-lines are theoretical since theses isotopes do not every occur in nature and can’t be constructed since they reject the extra nucleons instantaneously since they are so energetically unfavourable. On the proton side the drip-line defining this exists close to the “stable” on region on the edge of the valley of stability, since proton repulsion increases potential energy of adding new protons. On the neutron side the line is further from the region of stability since the extra repulsion doesn’t exist. At the edge of the region of stability some isotopes with long enough lifetimes to be observable which decay via proton and neutron emission exist. There are more proton cases since the proton repulsion makes this type of interaction more likely. Since the mechanism of decay in this case is not weak decay probabilities can be larger and lifetimes smaller.
Looking at stars
Emission and absorption (fingerprint) lines
The surface of the Sun contains all the elements found in the Solar System
The Origin of the Universe
Diagram of Evolution of the universe from the Big Bang (left) – to the present
During the dark ages no stars were created and after the first 3 minutes 98% of the Universe was composed of hydrogen and helium nuclei with a great deal of leftover energy.
In physical cosmology, Big Bang nucleosynthesis (abbreviated BBN, also known as primordial nucleosynthesis) refers to the production of nuclei other than those of the lightest isotope of hydrogen (hydrogen-1, 1H, having a single proton as a nucleus) during the early phases of the universe. Primordial nucleosynthesis is believed by most cosmologists to have taken place from 10 seconds to 20 minutes after the Big Bang, and is calculated to be responsible for the formation of most of the universe’s helium as the isotope helium-4 (4He), along with small amounts of the hydrogen isotope deuterium (2H or D), the helium isotope helium-3 (3He), and a very small amount of the lithium isotope lithium-7 (7Li). In addition to these stable nuclei, two unstable or radioactive isotopes were also produced: the heavy hydrogen isotope tritium (3H or T); and the beryllium isotope beryllium-7 (7Be); but these unstable isotopes later decayed into 3He and 7Li, as above.
Essentially all of the elements that are heavier than lithium and beryllium were created much later, by stellar nucleosynthesis in evolving and exploding stars.
The expansion of the Universe
1) The Big bang 2) Quark-gluon plasma at E-6s 3) Proton and neutron formation at E-4s 4) Formation of low mass nuclei at 3 minutes
5) Formation of neutral atoms at 400000 years 6) Star formation at E9 years 7) Dispersion of massive elements at greater than E9 years 8) Today approximately 13.772 billion years
The life cycle of Stars
Stage 1- Stars are born in a region of high density Nebula, and condenses into a huge globule of gas and dust and contracts under its own gravity.
Stage 2 – A region of condensing matter will begin to heat up and start to glow forming Protostars. If a protostar contains enough matter the central temperature reaches 15 million degrees Celsius
Stage 3 – At this temperature, nuclear reactions in which hydrogen fuses to form helium can start.
Stage 4 – The star begins to release energy, stopping it from contracting even more and causes it to shine. It is now a Main Sequence Star.
Stage 5 – A star of one solar mass remains in main sequence for about 10 billion years, until all of the hydrogen has fused to form helium.
Stage 6 – The helium core now starts to contract further and reactions begin to occur in a shell around the core.
Stage 7 – The core is hot enough for the helium to fuse to form carbon. The outer layers begin to expand, cool and shine less brightly. The expanding star is now called a Red Giant.
Stage 8 – The helium core runs out, and the outer layers drift of away from the core as a gaseous shell; this gas that surrounds the core is called a Planetary Nebula.
Stage 9 – The remaining core (thats 80% of the original star) is now in its final stages. The core becomes a White Dwarf the star eventually cools and dims. When it stops shining, the now dead star is called a Black Dwarf.
Massive stars have a mass 3x times that of the Sun. Some are 50x that of the Sun
Stage 1 – Massive stars evolve in a similar way to small stars until it reaches its main sequence stage (see small stars, stages 1-4). The stars shine steadily until the hydrogen has fused to form helium (it takes billions of years in a small star, but only millions in a massive star).
Stage 2 – The massive star then becomes a Red Supergiant and starts off with a helium core surrounded by a shell of cooling, expanding gas.
Stage 3 – In the next million years a series of nuclear reactions occur forming different elements in shells around the iron core.
Stage 4 – The core collapses in less than a second, causing an explosion called a Supernova, in which a shock wave blows of the outer layers of the star. (The actual supernova shines brighter than the entire galaxy for a short time).
Stage 5 – Sometimes the core survives the explosion. If the surviving core is between 1.5 – 3 solar masses it contracts to become a tiny, very dense Neutron Star. If the core is much greater than 3 solar masses, the core contracts to become a Black Hole.
Our Sun is not a massive star and it isn’t a first generation star.
We know the Sun is not a first generation star because it contains some heavier elements. These elements could not exist just based on nuclear fusion inside stars. Instead, some of this material must have come from the supernovae of other stars.
The sun contains roughly 70 percent hydrogen and 28 percent helium. It also contains roughly 1.5 percent carbon, nitrogen and oxygen. This element profile fits that of a main-sequence star of low mass. However, the missing 0.5 percent of the sun’s mass is made up of heavier elements. The sun is not massive enough to create these elements through its own nuclear fusion.
Hydrogen burning in stars
Protons collide all the time in our Sun’s core, but there is no bound state of two protons because there aren’t any neutrons to hold them together. Protons can only fuse if one of them undergoes beta plus decay to become a neutron at the moment of the collision. The neutron and the remaining proton fuse to form a deuterium nucleus, and this can react with another proton to form 3He. The beta plus decay is mediated by the weak force so it’s relatively slow process anyway, and the probability of the beta plus decay happening at just the right time is extremely low, which is why proton fusion is relatively slow in the Sun. It takes gazillions of proton-proton collisions to form a single deuterium nucleus. This explains why the probability of helium forming is so low.
The bottleneck in solar fusion is getting two hydrogen nuclei, i.e. two protons, to fuse together.
Why the process takes millions of years
There are two steps:
Hydrogen is found in nature but a neutron isn’t. On its own a neutron will decay within 10 minutes to a proton.
The energy required for the reaction is 1.8MeV. This cannot happen by applying classical physics but it can happen sometimes by applying quantum mechanics.
In quantum mechanics, the uncertainty principle, also known as Heisenberg’s uncertainty principle, is any of a variety of mathematical inequalities asserting a fundamental limit to the precision with which certain pairs of physical properties of a particle, known as complementary variables, such as position x and momentum p, can be known simultaneously. Introduced first in 1927, by the German physicist Werner Heisenberg, it states that the more precisely the position of some particle is determined, the less precisely its momentum can be known, and vice versa. The formal inequality relating the standard deviation of position σx and the standard deviation of momentum σp was derived by Earle Hesse Kennard later that year and by Hermann Weyl in 1928
(ħ is the reduced Planck constant, h / 2π)
We then have the bottle neck, then
This is never observed in a laboratory
Hydrogen burning takes millions of years
How are elements heavier than helium made?
The triple-alpha process is a set of nuclear fusion reactions by which three helium-4 nuclei (alpha particles) are transformed into carbon.
Older stars start to accumulate helium produced by the proton–proton chain reaction and the carbon–nitrogen–oxygen cycle in their cores. The products of further nuclear fusion reactions of helium with hydrogen or another helium nucleus produce lithium-5 and beryllium-8 respectively, both of which are highly unstable and decay almost instantly back into smaller nuclei. When the star starts to run out of hydrogen to fuse, the core of the star begins to collapse until the central temperature rises to 108 K (8.6 keV). At this point helium nuclei are fusing together faster than their product, beryllium-8, and decays back into two helium nuclei.
Once beryllium-8 is produced a little faster than it decays, the number of beryllium-8 nuclei in the stellar core increases to a large number. Then in its core there will be many beryllium-8 nuclei that can fuse with another helium nucleus to form carbon-12, which is stable:
The net energy release of the process is 1.166 pJ.
Because the triple-alpha process is unlikely, it needs a long time to produce much carbon. One consequence of this is that no significant amount of carbon was produced in the Big Bang because within minutes after the Big Bang, the temperature fell below the critical point for nuclear fusion.
Ordinarily, the probability of the triple alpha process is extremely small. However, the beryllium-8 ground state has almost exactly the energy of two alpha particles. In the second step, 8Be + 4He have almost exactly the energy of an excited state of 12C. These resonances greatly increase the probability that an incoming alpha particle will combine with beryllium-8 to form carbon. The existence of this resonance was predicted by Fred Hoyle before its actual observation, based on the physical necessity for it to exist, in order for carbon to be formed in stars. In turn, prediction and then discovery of this energy resonance and process gave very significant support to Hoyle’s hypothesis of stellar nucleosynthesis, which posited that all chemical elements had originally been formed from hydrogen, the true primordial substance.
The Hoyle State
The Hoyle state is an excited, spinless, resonant state of carbon-12. It is produced via the triple-alpha process, and was predicted to exist by Fred Hoyle in 1954. The existence of the Hoyle state is essential for the nucleosynthesis of carbon in helium-burning red giant stars, and predicts an amount of carbon production in a stellar environment which matches observations. The existence of the Hoyle state has been confirmed experimentally, but its precise properties are still being investigated. In 2011, an ab initio calculation of the low-lying states of carbon-12 found (in addition to the ground and excited spin-2 state) a resonance with all of the properties of the Hoyle state.
No Hoyle state means no elements beyond helium
The energy of the Hoyle state determines the amount of heavier elements
Sir Fred Hoyle FRS (24 June 1915 – 20 August 2001) was an English astronomer noted primarily for the theory of stellar nucleosynthesis and his often controversial stances on other scientific matters—in particular his rejection of the “Big Bang” theory, a term coined by him on BBC radio.
The Onion Structure of Heavy Stars
Elements up to Iron can be formed. It has the most stable nucleus.
Hydrogen burning takes about 1 x E9 years
Helium burning takes about 1 x E6 years
Carbon burning takes about 1 x E3 years
Neon and Oxygen burning takes about a year
Silicon burning takes about 1 x E-2 years
Creation of elements up to iron in burning (fusion) processes in stars is understood
How were the elements from iron to uranium made?
The abundances of the elements for A = 70 to 210
Double peaks at neutron numbers 46/50, 76/82, 116/126
Narrow peaks at N = 50, 82, 126 are magic numbers (equivalent to noble gases). Nuclei with magic number of neutrons are very stable against neutron capture.
The half-life of a neutron is 614 seconds and the decay has to happen in stars.
They are due to two separate processes and taken together they account for a majority of galactic chemical evolution of elements heavier than iron.
1) Slow neutron capture process: s process
Low neutron flux; beta decay < neutron capture time (this is well understood)
Abundance peaks at A = 84, 138, 208
The s-process or slow-neutron-capture-process is a nucleosynthesis process that occurs at relatively low neutron density and intermediate temperature conditions in stars. Under these conditions heavier nuclei are created by neutron capture, increasing the atomic weight of the nucleus by one. A neutron in the new nucleus decays by beta-minus decay to a proton, creating a nucleus of higher atomic number. The rate of neutron capture by atomic nuclei is slow relative to the rate of radioactive beta-minus decay, hence the name.
Magic numbers of neutrons in nuclei act as bottlenecks and are seen as s-process peaks.
The s process terminates in bismuth, the heaviest “stable” element, and polonium, the first non-primordial element after bismuth. (Bismuth is actually slightly radioactive, but with a half-life so long—a billion times the present age of the universe—that it is effectively stable over the lifetime of any existing star.)
The s-process acting in the range from Ag to Sb.
The process primarily occurs in AGB stars and is secondary, meaning that it requires pre-existing heavy isotopes as seed nuclei to be converted into other heavy nuclei.
2) Rapid neutron capture process: r process
High neutron flux
Beta-decay time > neutron capture time
Abundance peaks at A = 80, 130, 195
The r-process is a nucleosynthesis process that occurs in core-collapse supernovae and is responsible for the creation of approximately half of the neutron-rich atomic nuclei heavier than iron. The process entails a succession of rapid neutron captures (hence the name r-process) by heavy seed nuclei, typically 56Fe or other more neutron-rich heavy isotopes.
Most neutron-rich isotopes of elements heavier than nickel are produced, either exclusively or in part, by the beta decay of very radioactive matter synthesized during the r-process by rapid absorption, one after another, of free neutrons created during the explosions. The creation of free neutrons by electron capture during the rapid collapse to high density of the supernova core along with assembly of some neutron-rich seed nuclei makes the r-process a primary process; namely, one that can occur even in a star of pure H and He, in contrast to the B2FH designation as a secondary process building on pre-existing iron.
The r-process is responsible for our natural cohort of radioactive elements, such as uranium and thorium, as well as the most neutron-rich isotopes of each heavy element. In other words elements heavier than bismuth/polonium are produced this way.
The exact site of r-process is still unconfirmed however due to the conditions necessary (high neutron density, high temperature) core collapse supernovae and neutron star mergers are the most likely candidates.
The most probable candidate sites for the r-process has long been suggested to be core-collapse supernovae (spectral Type Ib, Ic and II), which may provide the necessary physical conditions for the r-process. However, the abundance of r-process nuclei requires that either only a small fraction of supernovae eject r-process nuclei to the interstellar medium, or that each supernova ejects only a very small amount of r-process material. In addition, the ejected material must be relatively neutron-rich, a condition which has been difficult to achieve in models. An alternative site proposed in 1974 was decompressing neutron star matter. It was proposed such matter is ejected from neutron stars merging with black holes in compact binaries. In 1989 this scenario was extended to binary neutron star mergers (a binary star system of two neutron stars that collide). These sites may now be starting to be observationally confirmed.
The r-process can only take place in extreme conditions such as when the core of a supernova collapses to become a neutron star, releasing tremendous energy in the process. Under those extreme neutron-rich conditions, atomic nuclei absorb neutrons to become increasingly heavy, and then undergo beta decay, leaving the nucleus one element higher on the periodic table. Gradually, the nuclei creep up the periodic table, leading to the creation of new elements.
Researchers at NAOJ and the University of Tokyo revealed that gold, platinum, and rare earth elements are likely produced and distributed through binary neutron star mergers. These elements are examples of a group known as r-process elements, which are produced by rapid neutron capture reactions. It has been a mystery where in nature the extreme conditions needed for this process could occur. The research team determined that r-process elements ejected from neutron star mergers at very high speeds (10 – 30 % the speed of light) quickly pervade their entire host galaxy, and can account for the chemical signature observed in stars in the Milky Way Galaxy and nearby dwarf galaxies.
The r-processes are believed to occur during rare core-collapse supernovae and neutron star mergers. More information can be found at
The mystery of the r-process
It requires high neutron densities and temperatures and these can be found in supernovae explosions and neutron star mergers.
What are the properties of the nuclei involved? The problem is that the majority of them have never been studied.
Description: This movie demonstrates nucleosynthesis of heavy elements in neutron-rich ejecta from ns/bh encounters. The axis represents proton and neutron numbers, and each cell is colour-coded depending on the abundance of a particular isotope. Insets demonstrate the dynamics of the integrated abundances and the evolution of thermodynamic parameters in the ejected matter.
Nucleosynthesis in r-process
In the animations, the term encounterd refers to both gravitational wave-driven binary mergers, and dynamical collisions which may occur in dense environments such as globular clusters. Much interesting work has recently been done in this direction, for a review and a pointer to literature see, for example,
Compact binary mergers: an astrophysical perspective S. Rosswog, (2011), review “Nuclei in the Cosmos 2010”, eprint arXiv:1012.0912.
Paper II: The Electromagnetic Signals of Compact Binary Mergers
T. Piran, E. Nakar and S. Rosswog, submitted to MNRAS (2012), eprint:arXiv:1204.6242
Merger of two neutron stars, no spin: 3D cut of the temperature
Description: Animation of the coalescence of two neutron stars with masses 1.4 and 1.3 M☉, with no spin (referred to as “run H” in Paper I, see Table 1). Considered the most likely case of a neutron star merger.
- Temperature, 3D-rendering, annotated: .mov (23 Mb)
- Temperature, 3D-rendering, no annotation: .mov (29 Mb)
- Temperature, 3D-cut, no annotation: .mov (4.2 Mb)
- Electron fraction (Ye), 3D-cut, annotated: .mov (14 Mb)
Description: Coalescence of two corotating neutron stars with masses 1.4 and 1.3 M☉ (referred to as “run G” in Paper I, see Table 1). This case is somewhat academic, since the neutron star viscosity is far too low to lead to a tidal locking during neutron star inspiral.
Merger of two corotating neutron stars: 3D-cut of the temperature
- Temperature, 3D-cut, annotated: .mov (23 Mb)
- Temperature, 3D-rendering, no annotation: .mov (29 Mb)
- Temperature, 3D-rendering, annotated: .mov (26 Mb)
- Density, 3D-cut, annotated: .mov (3.5 Mb)
For collisions presented below, only parabolic (marginally unbound) cases were considered. The strength of a collision is characterized by the quantity β≡ (r1 + r2) /rperi, where the object radii ri refer in the black hole cases to the Schwarzschild radii. See Paper I for more information.
Neutron star – neutron star collisions
For the double neutron star (ns2)-collisions presented below we fix a slight mass asymetry (m1= 1.3 M☉, m1= 1.4 M☉) and explore the dependence on the impact strength parameter β.
Description: Animation of a neutron star – neutron star grazing impact. Parameters: m1= 1.3 M☉, m1= 1.4 M☉, β=1. This simulation is referred to as “run A” in Paper I (see Table 1).
Link to download the movie: .mov (36 Mb)
Grazing impact of two neutron stars with impact parameter β=1. Movie shows 3D-cuts of the temperature distribution
Description: Animation of a neutron star – neutron star grazing impact. Parameters: m1= 1.3 M☉, m1= 1.4 M☉, β=2. This simulation is referred to as “run B” in Paper I (see Table 1).
Link to download the movie: .mov (15 Mb)
More central impact of two neutron stars with impact parameter β=2
Description: Neutron star – neutron star grazing impact. Parameters: m1= 1.3 M☉, m1= 1.4 M☉, β=5. This simulation is referred to as “run C” in Paper I (see Table 1).
Link to download the movie: .mov (17 Mb)
Very central impact of two neutron stars; Impact parameter β=5
1) Excite a nucleus
3) Excite the nucleus
Then gamma rays (bursts of energy)
Atomic mass evaluations – A good description of what is known
FAIR (Facility for Anti-proton and Ion Research)
FAIR is a new, unique international accelerator facility for the research with antiprotons and ions. It is ready to be built within the coming years near Darmstadt in Hesse, Germany. The major part of the budget will be provided by the Federal Republic of Germany, together with the State of Hesse. Other fractions will be funded by international partners from Europe and overseas.
The new facility, where various physics programs can be operated in parallel, will offer outstanding research opportunities and discovery potential for about 3000 scientists from about 50 countries. In the course of the coming decades the experiments will reveal consolidated findings about so far unknown states of matter and still missing information about the evolution of the Universe 13.8 billion years ago.
Solar System abundances
Universality of the r-process (there are other stars with the same abundances)
Observational nuclear astrophysics: neutron-capture element abundances in old, metal-poor stars
The chemical abundances of metal-poor stars provide a great deal of information regarding the individual nucleosynthetic processes that created the observed elements and the overall process of chemical enrichment of the galaxy since the formation of the first stars. The paper reviews the abundance patterns of the neutron-capture elements (Z ≥ 38) in those metal-poor stars and our current understanding of the conditions and sites of their production at early times. It reviews the relative contributions of these different processes to the build-up of these elements within the galaxy over time, and outline outstanding questions and uncertainties that complicate the interpretation of the abundance patterns observed in metal-poor stars. It is anticipated that future observations of large samples of metal-poor stars will help discriminate between different proposed neutron-capture element production sites and better trace the chemical evolution of the galaxy.
The origin of the elements
Note: yellow-red all related to massive stars (>8-12 solar masses)
Elements are created in nuclear reactions
Fusion processes in heavy stars create elements up to iron
Proton + proton joining takes billions of years
All nuclei heavier than helium created via the Hoyle state
Elements beyond iron created in a neutron rich environment
We still do not know where (although it isn’t in a lab). Is it neutron star mergers and/or supernovae? How are heavy elements made? Is it the r-process?