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science teacher 2010 Featuring: Iron Weathering steel Black sand Iron in the stars Iron and origin of life Iron fertilisation Dietary iron Iron, oxygen, and life Plus: Ripping yarns Science writing and the media What is ‘Western’ about science And more... Number 124 ISSN 0110-7801 NZ iron in stars science teacher 124 Iron is only one of almost 100 elements of the periodic table that are found in the Sun and other stars, but it is one of the most important elements in how it affects the evolution of a star. Iron is the last element that is fused at the core of the star before it disintegrates in a supernova. Since no more energy can be gained to hold up the star, the star undergoes gravitational collapse. The reason no more energy can be gained comes down to how the fundamental components of the nucleus combine together. How is the nucleus held together? If you measure the mass of an atomic nucleus and then add up the mass of the all its constituent protons and neutrons, you will find that they do not have the same value. The nucleus has less mass than the parts that make it up. If you convert this mass difference to an energy using E=mc2, this value is called the binding energy. It can be thought of as a measure of how much energy is required to break apart the nucleus. Dividing this binding energy by the number of nucleons (the number of protons plus the number of neutrons) gives the binding energy per nucleon. This number is the key reason why iron is so important to the life of stars (see Figure 1). Average binding energy per nucleon (MeV) 9 8 7 6 5 4 3 2 1 0 0 30 60 90 120 150 180 210 240 270 Number of nucleons in nucleus Figure 1: As you change the number of nucleons in a nucleus, there is also a change in the binding energy per nucleon. Due to the interplay between the strong nuclear force and electromagnetic force there is a peak in this energy around iron-56. As you move through the periodic table from hydrogen to sodium (the stable isotope of which has a nucleus containing 23 nucleons – 11 protons and 12 neutrons), there is an increase in the binding energy per nucleon. As more nucleons are added to the nucleus, the combined strong nuclear force packs the nucleus tighter, making it harder to break apart. The strong nuclear force is a fundamental force that keeps the positively charged protons in the nucleus together. From magnesium (24 nucleons) to xenon (131 nucleons) there is a relatively flat section. This feature is associated with the growing importance of the electromagnetic repulsion between the protons. The strong nuclear force and the electromagnetic force are able to keep each other in balance, such that there is no real change in the energy needed to break apart the nucleus. Beyond caesium (134 nucleons), there is a slow decline as the nucleons on one side of the nucleus cannot exert a strong enough nuclear force on the nucleons on the other side but the electromagnetic repulsion is being felt. It is this curve that tells us why a fission reactor uses uranium and a fusion power plant would use hydrogen. To the left of the peak, atoms become more tightly bound if they are combined to form a single nucleus that lies nearer to the peak. This process is called fusion. For atoms to the right of the peak, they would be more tightly bound by breaking into atoms that are nearer to the peak. This process is fission. It is during the flat section from magnesium to xenon that there is a peak at around iron. The peak actually falls at nickel-62, not iron-56 as is commonly stated. In fact, iron-58 also has a higher binding energy per nucleon (Fewell). The reason why iron-56 is more important than these elements to the life cycle of some stars comes down to helium nucleus. iron Iron is the last element that is fused at the core of the star before it disintegrates in a supernova, as Jeffrey Simpson, post graduate student in the Department of Physics and Astronomy at the University of Canterbury, explains: How do stars evolve and explode? Stars shine as a result of the release of energy due to nuclear fusion. Their evolution is primarily controlled by their mass. In a massive star (5 times the mass of the Sun or more), the evolution will follow a different track to that of the Sun. This is due to its size which allows for higher energies and pressures to be obtained at its core where the fusion takes place. After it has exhausted its fuel of hydrogen at the core, by fusing it to helium, the core contracts under gravitational pressure. This pulls hydrogen down to deeper layers so that a shell of hydrogen fusion forms around the nonburning helium core. This shell of hydrogen then causes the outer layers to expand due to the increased thermal energy, producing a red giant. Eventually the pressure and temperature at the core reaches a level that helium fusion can take through what is known as the triplealpha process. Three helium nuclei are fused to form one carbon-12. At a temperature of 200 million kelvin, a carbon atom can fuse with another helium nuclei or alpha particle to form oxygen-16. As the star ages, higher temperatures are reached as one source of fuel is used up and the star’s core contracts. Over time the conditions are reached such that more alpha particles can be captured to form neon-20, silicon-28, and then finally nickel-56 and iron-56 (see Figure 2). The star is now like an onion (see Figure 3), with layers of different materials tracing its chemical evolution. At its core is iron and nickel. As stated above, nickel-62 is New Zealand Association of Science Educators 21 NZ science teacher iron 124 Figure 2: Carbon burning. Two carbon atoms fuse to form a neon-20 nucleus plus a helium nucleus and some energy in the form of a gamma ray. Diagram adapted by Jeffrey Simpson from: http://en.wikipedia.org/wiki/File:CNO_Cycle.svg actually the most stable element in terms of its nuclear binding energy per nucleon. But it is difficult to create in the interior of stars through the addition of abundant species. It is almost certainly produced in stars, but not in great enough quantities to compete with iron in terms of being relevant to stellar evolution. The iron/nickel core mass increases until it reaches the Chandrasekhar limit, which is about 1.4 times the mass of the Sun. At this mass, electron degeneracy pressure is overcome. This is the pressure caused by the Pauli Exclusion Principle which forbids fermions (such as electrons) being in the same energy state. It is an extremely strong force but can be overcome. At this point, electrons and protons combine to form neutrons and the core collapses, leading to a supernova. There are several different types of supernovae. In terms of our story of iron in astronomy, it is the type Ib, Ic and Figure 3: This diagram shows a simplified (not to scale) cross-section of a massive, evolved star (with a mass greater than eight times the Sun). Where the pressure and temperature permit, concentric shells of Hydrogen (H), Helium (He), Carbon (C), Neon/Magnesium (Ne), Oxygen (O) and Silicon (Si) plasma are burning inside the star. Ref: http://commons.wikimedia.org/wiki/File:Evolved_star_fusion_shells.svg 22 New Zealand Association of Science Educators II supernovae that we are interested in. These are the supernovae where iron plays a key role. So iron having nearly the highest binding energy per nucleon is what causes a core collapse supernova such as the one that caused SN 1054, creating a source of light so bright it could be seen during the day. This was the progenitor of the Crab Nebula (see Figure 4). Another example of a core-collapse supernova was SN 1987A, which occurred in the Large Magellanic Cloud (Figure 5). One of its discoverers was a New Zealander, amateur astronomer Albert Jones who has made 500,000 observations of variable stars. The other type of supernova is Type 1a (Figure 6). This is caused by a white dwarf star exceeding the Chandrasekhar limit. White dwarfs are the remnants of stars like our Sun which do not explode in supernova. Instead, they do not have any more fusion at their core beyond carbon and oxygen. Their outer layers are lost and the core is left naked in space. In type 1a supernova, the white dwarf gains matter from a companion star which allows it to exceed the Chandrasekhar limit. Here iron is produced, but this is explosively in the core of the star. It is also produced by the radioactive decay of nickel-56 to cobalt-56 and then iron-56. Figure 4: The Crab Nebula as seen by the Hubble Space Telescope in 2000. Japanese and Chinese astronomers recorded this violent event nearly 1,000 years ago in 1054. It is believed to be the result of a type II supernova. Ref: http://hubblesite.org/newscenter/archive/releases/2005/37/) NZ science teacher 124 iron Figure 5: This Hubble Space Telescope image, taken in February 1994, with the Wide Field and Planetary Camera 2, shows the full system of three rings of glowing gas surrounding supernova 1987A. Courtesy of P. Challis (Harvard-Smithsonian Center for Astrophysics Are there stars with no iron or other heavy elements in their atmosphere? So far, no star has been observed that has no iron or other heavy elements in its atmosphere. But the current theories of Big Bang cosmology predict there must have been some stars which were composed almost only of hydrogen and helium. These are known as population III stars. A few minutes after the Big Bang there was a period of primordial nucleosynthesis: the creation of elements from the sea of protons and neutrons that filled the Universe. This resulted in the production of deuterium (a hydrogen isotope of one proton and one neutron), helium-3 and -4, and lithium-6 and -7. It is from these elements that the first stars would have formed. These population III stars are predicted to be much more massive than stars that exist today. Model simulations predict that they would have been over 100 times the mass of the Sun. Due to their massive size they would have lived for only a couple of million years before exploding in a supernova. For stars in the mass range of 130 solar masses, this would result in a supernova for which is predicted a large proportion of the heavy elements created would take the form of iron. Population III stars have not been observed yet. This is due to the extremely short lifetimes and existing only in the very early Universe. The most metal-poor stars that are observed today are found in groups of stars called globular clusters. These are spherical conglomerations of between 10,000 and a million stars that form halos around galaxies. One of the stars with the lowest known amount of heavy elements is HE0107-5240, with about 1/200,000 of the iron content of the Sun (Lau, Stancliffe and Tout). What about the other extreme? That is, stars that have an Figure 6: The spiral galaxy NGC 2770 and its two supernovae. The bright star at the edge of the galaxy in the top right is SN 2008D, while the star left of the centre is SN 2007uy. Ref: http://www.eso.org/public/images/eso0823a/ equivalently high (200,000 times) abundances of iron? No such stars are known, though it is theorized by some that the surfaces of neutron stars could consist of iron. Neutron stars are the remnant of supernova. They consist of neutrons supported by the Pauli Exclusion Principle. But their surface regions are thought to be composed of atomic nuclei, with the possibility of it being iron left from the core of the massive star that underwent the supernova. Conclusion Through the interplay of the strong nuclear force and electromagnetic force, iron finds itself with one of the highest binding energies per nucleon. In addition, iron56 can be created by the addition of alpha particles in the interior of stars allowing it to be built up in the core of massive stars. It is here that iron acts as one of the final fusion products. With no more energy available, the star’s core collapses, causing a supernova. For further information contact: [email protected] Bibliography Fewell, M. P. (1995). “The atomic nuclide with the highest mean binding energy.” American Journal of Physics, 63.7, 653-658. Fleurot, F. (2010). Evolution of Massive Stars. 20 April 2010, http://nu.phys. laurentian.ca/~fleurot/evolution/ Halliday, D., Resnick, R., & Walker, J. (2005). Fundamentals of Physics. 7th Edition. Wiley. Lau, Herbert H.B., Stancliffe, R. J., & Tout, C. A. (2007). “Carbon-rich extremely metal poor stars: signatures of Population III asymptotic giant branch stars in binary systems.” Monthly Notices of the Royal Astronomical Society, 378.2, 563-568. New Zealand Association of Science Educators 23