Nuclear Binding energy

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Nuclear binding energy is the energy required to split the nucleus of an atom into its component parts. The component parts are neutrons and protons, which are collectively called nucleons. The binding energy of nuclei is usually a positive number, since most nuclei require net energy to separate them into individual protons and neutrons. Thus, the mass of an atom's nucleus is usually less than the sum of the individual masses of the constituent protons and neutrons when separated. This notable difference is a measure of the nuclear binding energy, which is a result of forces that hold the nucleus together. During the splitting of the nucleus, some of the mass of the nucleus (i.e. some nucleons) gets converted into huge amounts of energy (according to Einstein's equation E=mc²) and thus this mass is removed from the total mass of the original particles, and the mass is missing in the resulting nucleus. This missing mass is known as the mass defect, and represents the energy released when the nucleus is formed.
The term nuclear binding energy may also refer to the energy balance in processes in which the nucleus splits into fragments composed of more than one nucleon, and in this case the binding energies for the fragments, as compared to the whole, will be higher. If new binding energy is available when light nuclei fuse, or when heavy nuclei split, either of these processes result in releases of the binding energy. This energy, available as nuclear energy, can be used to produce electricity (nuclear power) or as a nuclear weapon. When a large nucleus splits into pieces, excess energy is emitted as photons (gamma rays) and as kinetic energy of a number of different ejected particles (nuclear fission products).

The nuclear binding energies and forces are on the order of a million times greater than the electron binding energies of light atoms like hydrogen.^[1]

The mass defect of a nucleus represents the mass of the energy of binding of the nucleus, and is the difference between the mass of a nucleus and the sum of the masses of the nucleons of which it is composed. Determining the relevant nuclear binding energy encompasses three steps of calculation, which involves the creation of mass defect by removing the mass as released energy.^[2]

Introduction

Binding energy per nucleon of common isotopes.

Nuclear binding energy is explained by the basic principles involved in nuclear physics.

Nuclear energy

An absorption or release of nuclear energy occurs in nuclear reactions or radioactive decay; those that absorb energy are called endothermic reactions and those that release energy are exothermic reactions. Energy is consumed or liberated because of differences in the nuclear binding energy between the incoming and outgoing products of the nuclear transmutation.^[3]

The best-known classes of exothermic nuclear transmutations are fission and fusion. Nuclear energy may be liberated by atomic fission, when heavy atomic nuclei (like uranium and plutonium) are broken apart into lighter nuclei. The energy from fission is used to generate electric power in hundreds of locations worldwide. Nuclear energy is also released during atomic fusion, when light nuclei like hydrogen are combined to form heavier nuclei such as helium. The Sun and other stars use nuclear fusion to generate thermal energy which is later radiated from the surface, a type of stellar nucleosynthesis. In any exothermic nuclear process, nuclear mass might ultimately be converted to thermal energy, given off as heat, carries away the mass with it.

In order to quantify the energy released or absorbed in any nuclear transmutation, one must know the nuclear binding energies of the nuclear components involved in the transmutation.

The nuclear force

Electrons and nuclei are kept together by electrostatic attraction (negative attracts positive). Furthermore, electrons are sometimes shared by neighboring atoms or transferred to them (by processes of quantum physics), and this link between atoms is referred to as a chemical bond, and is responsible for the formation of all chemical compounds.^[4]

The force of electric attraction does not hold nuclei together, because all protons carry a positive charge and repel each other. Thus, electric forces do not hold nuclei together, because they act in the opposite direction. It has been established that binding neutrons to nuclei clearly requires a non-electrical attraction.^[4]

Therefore, another force, called the nuclear force (or residual strong force) holds the nucleons of nuclei together. This force is a residuum of the strong interaction, which binds quarks into nucleons at an even smaller level of distance.

The nuclear force must be stronger than the electric repulsion at short distances, but weaker far away, or else different nuclei might tend to clump together. Therefore it has short-range characteristics. An analogy to the nuclear force is the force between two small magnets: magnets are very difficult to separate when stuck together, but once pulled a short distance apart, the force between them drops almost to zero.^[4]

Unlike gravity or electrical forces, the nuclear force is effective only at very short distances. At greater distances, the electrostatic force dominates: the protons repel each other because they are positively charged, and like charges repel. For that reason, the protons forming the nuclei of ordinary hydrogen—for instance, in a balloon filled with hydrogen—do not combine to form helium (a process that also would require some to combine with electrons and become neutrons). They cannot get close enough for the nuclear force, which attracts them to each other, to become important. Only under conditions of extreme pressure and temperature (for example, within the core of a star), can such a process take place.^[5]

Physics of nuclei

The nuclei of atoms are found in many different sizes. In hydrogen they contain just one proton, in deuterium or heavy hydrogen a proton and a neutron; in helium, two protons and two neutrons, and in carbon, nitrogen and oxygen - six, seven and eight of each particle, respectively. A helium nucleus weighs less than the sum of the weights of its components. The same phenomenon is found for carbon, nitrogen and oxygen. For example, the carbon nucleus is slightly lighter than three helium nuclei, which can combine to make a carbon nucleus. This illustrates the mass defect.

Mass defect

The fundamental reason for the "mass defect" is Albert Einstein's formula E = m c², expressing the equivalence of energy and mass. By this formula, adding energy also increases mass (both weight and inertia), whereas removing energy decreases mass.

If a combination of particles contains extra energy—for instance, in a molecule of the explosive TNT—weighing it reveals some extra mass, compared to its end products after an explosion. (The weighing must be done after the products have been stopped and cooled, however, as the extra mass must escape from the system as heat before its loss can be noticed, in theory.) On the other hand, if one must inject energy to separate a system of particles into its components, then the initial weight is less than that of the components after they are separated. In the latter case, the energy injected is "stored" as potential energy, which shows as the increased mass of the components that store it. This is an example of the fact that energy of all types is seen in systems as mass, since mass and energy are equivalent, and each is a "property" of the other.

The latter scenario is the case with nuclei such as helium: to break them up into protons and neutrons, one must inject energy. On the other hand, if a process existed going in the opposite direction, by which hydrogen atoms could be combined to form helium, then energy would be released. The energy can be computed using E = Δm c² for each nucleus, where Δm is the difference between the mass of the helium nucleus and the mass of four protons (plus two electrons, absorbed to create the neutrons of helium).

For elements heavier than oxygen, the energy that can be released by assembling them from lighter elements decreases, up to iron. For nuclei heavier than iron, one actually releases energy by breaking them up into 2 fragments. That is how energy is extracted by breaking up uranium nuclei in nuclear power reactors.

The reason the trend reverses after iron is the growing positive charge of the nuclei. The electric force may be weaker than the nuclear force, but its range is greater: in an iron nucleus, each proton repels the other 25 protons, while the nuclear force only binds close neighbors.

As nuclei grow bigger still, this disruptive effect becomes steadily more significant. By the time polonium is reached (84 protons), nuclei can no longer accommodate their large positive charge, but emit their excess protons quite rapidly in the process of alpha radioactivity—the emission of helium nuclei, each containing two protons and two neutrons. (Helium nuclei are an especially stable combination.) Because of this process, nuclei with more than 98 protons are not found naturally on Earth. The isotopes beyond uranium (atomic number 92) with the longest half-lives are plutonium-244 (80 million years) and curium-247 (16 million years).

Solar binding energy

The nuclear fusion process works as follows: five billion years ago, the new Sun formed when gravity pulled together a vast cloud of gas and dust, from which the Earth and other planets also arose. The gravitational pull released energy and heated the early Sun, much in the way Helmholtz proposed.

Thermal energy appears as the motion of atoms and molecules: the higher the temperature of a collection of particles, the greater is their velocity and the more violent are their collisions. When the temperature at the center of the newly formed Sun became great enough for collisions between nuclei to overcome their electric repulsion, and bring them into the short range of the attractive nuclear force, nuclei began to stick together. When this began to happen, protons combined into deuterium and then helium, with some protons changing in the process to neutrons (plus positrons, positive electrons, which combine with electrons and are destroyed). This released nuclear energy now keeps up the high temperature of the Sun's core, and the heat also keeps the gas pressure high, keeping the Sun at its present size, and stopping gravity from compressing it any more. There is now a stable balance between gravity and pressure.

Different nuclear reactions may predominate at different stages of the Sun's existence, including the proton-proton reaction and the carbon-nitrogen cycle—which involves heavier nuclei, but whose final product is still the combination of protons to form helium.

A branch of physics, the study of controlled nuclear fusion, has tried since the 1950s to derive useful power from nuclear fusion reactions that combine small nuclei into bigger ones, typically to heat boilers, whose steam could turn turbines and produce electricity. Unfortunately, no earthly laboratory can match one feature of the solar powerhouse: the great mass of the Sun, whose weight keeps the hot plasma compressed and confines the nuclear furnace to the Sun's core. Instead, physicists use strong magnetic fields to confine the plasma, and for fuel they use heavy forms of hydrogen, which burn more easily. Magnetic traps can be rather unstable, and any plasma hot enough and dense enough to undergo nuclear fusion tends to slip out of them after a short time. Even with ingenious tricks, the confinement in most cases lasts only a small fraction of a second.

Combining nuclei

Small nuclei that are larger than hydrogen can combine into bigger ones and release energy, but in combining such nuclei, the amount of energy released is much smaller compared to hydrogen fusion. The reason is that while the overall process releases energy from letting the nuclear attraction do its work, energy must first be injected to force together positively charged protons, which also repel each other with their electric charge.^[5]

For elements that weigh more than iron (a nucleus with 26 protons), the fusion process no longer releases energy. In even heavier nuclei energy is consumed, not released, by combining similar sized nuclei. With such large nuclei, overcoming the electric repulsion (which affects all protons in the nucleus) requires more energy than what is released by the nuclear attraction (which is effective mainly between close neighbors). Conversely, energy could actually be released by breaking apart nuclei heavier than iron.^[5]

With the nuclei of elements heavier than lead, the electric repulsion is so strong that some of them spontaneously eject positive fragments, usually nuclei of helium that form very stable combinations (alpha particles). This spontaneous break-up is one of the forms of radioactivity behavior exhibited by some nuclei.^[5]

Nuclei heavier than lead (except for bismuth, thorium, uranium, and plutonium) spontaneously break up too quickly to appear in nature as primordial elements, though they can be produced artificially or as intermediates in the decay chains of lighter elements. Generally, the heavier the nuclei are, the faster they spontaneously decay.^[5]

Iron nuclei are the most stable nuclei (in particular iron-56), and the best sources of energy are therefore nuclei whose weights are as far removed from iron as possible. One can combine the lightest ones—nuclei of hydrogen (protons)—to form nuclei of helium, and that is how the Sun generates its energy. Or else one can break up the heaviest ones—nuclei of uranium or plutonium—into smaller fragments, and that is what nuclear power reactors do.^[5]

Nuclear binding energy

An example that illustrates nuclear binding energy is the nucleus of ¹²C (Carbon 12), which contains 6 protons and 6 neutrons. The protons are all positively charged and repel each other, but the nuclear force overcomes the repulsion and causes them to stick together. The nuclear force is a close-range force (it is very strongly inversely proportionate to distance), and virtually no effect of this force is observed outside the nucleus. The nuclear force also pulls neutrons together, or neutrons and protons.^[6]

The energy of the nucleus is negative with regard to the energy of the particles pulled apart to infinite distance (just like the gravitational energy of planets of the solar system), because energy must be utilized to split a nucleus into its individual protons and neutrons. Mass spectrometers have measured the masses of nuclei, which are always less than the sum of the masses of protons and neutrons that form them, and the difference—by the formula E = m c²—gives the binding energy of the nucleus.^[6]

Nuclear fusion

The binding energy of helium is the energy source of the Sun and of most stars. The sun is composed of 74 percent hydrogen (measured by mass), an element whose nucleus is a single proton. Energy is released in the sun when 4 protons combine into a helium nucleus, a process in which two of them are also converted to neutrons.^[6]

The conversion of protons to neutrons is the result of another nuclear force, known as the weak (nuclear) force. The weak force, like the strong force, has a short range, but is much weaker than the strong force. The weak force tries to make the number of neutrons and protons into the most energetically stable configuration. For nuclei containing less than 40 particles, these numbers are usually about equal. Protons and neutrons are closely related and are sometimes collectively known as nucleons. As the number of particles increases toward a maximum of about 209, the number of neutrons to maintain stability begins to outstrip the number of protons, until the ratio of neutrons to protons is about three to two.^[6]

The protons of hydrogen combine to helium only if they have enough velocity to overcome each other's mutual repulsion sufficiently to get within range of the strong nuclear attraction. This means that fusion only occurs within a very hot gas. Hydrogen hot enough for combining to helium requires an enormous pressure to keep it confined, but suitable conditions exist in the central regions of the Sun, where such pressure is provided by the enormous weight of the layers above the core, pressed inwards by the Sun's strong gravity. The process of combining protons to form helium is an example of nuclear fusion.^[6]

The earth's oceans contain a large amount of hydrogen that could theoretically be used for fusion, and helium byproduct of fusion does not harm the environment, so some consider nuclear fusion a good alternative to supply humanity's energy needs. Experiments to generate electricity from fusion have so far have only partially succeeded. Sufficiently hot hydrogen must be ionized and confined. One technique is to use very strong magnetic fields, because charged particles (like those trapped in the Earth's radiation belt) are guided by magnetic field lines. Fusion experiments also rely on heavy hydrogen, which fuses more easily, and gas densities can be moderate. But even with these techniques far more net energy is consumed by the fusion experiments than is yielded by the process.^[6]

The binding energy maximum and ways to approach it by decay

In the main isotopes of light nuclei, such as carbon, nitrogen and oxygen, the most stable combination of neutrons and of protons are when the numbers are equal (this continues to element 20, calcium). However, in heavier nuclei, the disruptive energy of protons increases, since they are confined to a tiny volume and repel each other. The energy of the strong force holding the nucleus together also increases, but at a slower rate, as if inside the nucleus, only nucleons close to each other are tightly bound, not ones more widely separated.^[6]

The net binding energy of a nucleus is that of the nuclear attraction, minus the disruptive energy of the electric force. As nuclei get heavier than helium, their net binding energy per nucleon (deduced from the difference in mass between the nucleus and the sum of masses of component nucleons) grows more and more slowly, reaching its peak at iron. As nucleons are added, the total nuclear binding energy always increases—but the total disruptive energy of electric forces (positive protons repelling other protons) also increases, and past iron, the second increase outweighs the first. Iron-56 (⁵⁶Fe) is the most efficiently bound nucleus^[6] meaning that it has the least average mass per nucleon. However, nickel-62 is the most tightly bound nucleus in terms of energy of binding per nucleon. (Nickel-62's higher energy of binding does not translate to a larger mean mass loss than Fe-56, because Ni-62 has a slightly higher ratio of neutrons/protons than does iron-56, and the presence of the heavier neutrons increases nickel-62's average mass per nucleon).

To reduce the disruptive energy, the weak interaction allows the number of neutrons to exceed that of protons—for instance, the main isotope of iron has 26 protons and 30 neutrons. Isotopes also exist where the number of neutrons differs from the most stable number for that number of nucleons. If the ratio of protons to neutrons is too far from stability, nucleons may spontaneously change from proton to neutron, or neutron to proton.

The two methods for this conversion are mediated by the weak force, and involve types of beta decay. In the simplest beta decay, neutrons are converted to protons by emitting a negative electron and an antineutrino. This is always possible outside a nucleus because neutrons are more massive than protons by an equivalent of about 2.5 electrons. In the opposite process, which only happens within a nucleus, and not to free particles, a proton may become a neutron by ejecting a positron. This is permitted if enough energy is available between parent and daughter nuclides to do this (the required energy difference is equal to 1.022 MeV, which is the mass of 2 electrons). If the mass difference between parent and daughter is less than this, a proton-rich nucleus may still convert protons to neutrons by the process of electron capture, in which a proton simply captures one of the atom's K orbital electrons, emits a neutrino, and becomes a neutron.^[6]

Among the heaviest nuclei, starting with tellurium nuclei (element 52) containing 106 or more nucleons, electric forces may be so destabilizing that entire chunks of the nucleus may be ejected, usually as alpha particles, which consist of two protons and two neutrons (alpha particles are fast helium nuclei). (Beryllium-8 also decays, very quickly, into two alpha particles.) Alpha particles are extremely stable. This type of decay becomes more and more probable as elements rise in atomic weight past 106.

The curve of binding energy is a graph that plots the binding energy per nucleon against atomic mass. This curve has its main peak at iron and nickel and then slowly decreases again, and also a narrow isolated peak at helium, which as noted is very stable. The heaviest nuclei in nature, uranium ²³⁸U, are unstable, but having a lifetime of 4.5 billion years, close to the age of the Earth, they are still relatively abundant; they (and other nuclei heavier than iron) may have formed in a supernova explosion ^[7] preceding the formation of the solar system. The most common isotope of thorium, ²³²Th, also undergoes α particle emission, and its half-life (time over which half a number of atoms decays) is even longer, by several times. In each of these, radioactive decay produces daughter isotopes that are also unstable, starting a chain of decays that ends in some stable isotope of lead.^[6]

Determining nuclear binding energy

Calculation can be employed to determine the nuclear binding energy of nuclei. The calculation involves determining the mass defect, converting it into energy, and expressing the result as energy per mole of atoms, or as energy per nucleon.^[2]

Conversion of mass defect into energy

Mass defect is defined as the difference between the mass of a nucleus, and the sum of the masses of the nucleons of which it is composed. The mass defect is determined by calculating three quantities.^[2] These are: the actual mass of the nucleus, the composition of the nucleus (number of protons and of neutrons), and the masses of a proton and of a neutron. This is then followed by converting the mass defect into energy. This quantity is the nuclear binding energy, however it must be expressed as energy per mole of atoms or as energy per nucleon.^[2]

Fission and fusion

Nuclear energy is released by the splitting (fission) or merging (fusion) of the nuclei of atom(s). The conversion of nuclear mass-energy to a form of energy, which can remove some mass when the energy is removed, is consistent with the mass-energy equivalence formula ΔE = Δm c², in which ΔE = energy release, Δm = mass defect, and c = the speed of light in a vacuum (a physical constant).

Nuclear energy was first discovered by French physicist Henri Becquerel in 1896, when he found that photographic plates stored in the dark near uranium were blackened like X-ray plates (X-rays had recently been discovered in 1895).^[8]

Nuclear chemistry can be used as a form of alchemy to turn lead into gold or change any atom to any other atom (though this may require many intermediate steps).^[7] Radionuclide (radioisotope) production often involves irradiation of another isotope (or more precisely a nuclide), with alpha particles, beta particles, or gamma rays. Nickel-62 has the highest binding energy per nucleon of any isotope. If an atom of lower average binding energy is changed into two atoms of higher average binding energy, energy is given off. Also, if two atoms of lower average binding energy fuse into an atom of higher average binding energy, energy is given off. The chart shows that fusion of hydrogen, the combination to form heavier atoms, releases energy, as does fission of uranium, the breaking up of a larger nucleus into smaller parts. Stability varies between isotopes: the isotope U-235 is much less stable than the more common U-238.

Nuclear energy is released by three exoenergetic (or exothermic) processes:

• Radioactive decay, where a neutron or proton in the radioactive nucleus decays spontaneously by emitting either particles, electromagnetic radiation (gamma rays), or both. Note that for radioactive decay, it is not strictly necessary for the binding energy to increase. What is strictly necessary is that the mass decrease. If a neutron turns into a proton and the energy of the decay is less than 0.782343 MeV (such as rubidium-87 decaying to strontium-87), the average binding energy per nucleon will actually decrease.
• Fusion, two atomic nuclei fuse together to form a heavier nucleus
• Fission, the breaking of a heavy nucleus into two (or more rarely three) lighter nuclei

Binding energy for atoms

The binding energy of an atom (including its electrons) is not the same as the binding energy of the atom's nucleus. The measured mass deficits of isotopes are always listed as mass deficits of the neutral atoms of that isotope, and mostly in MeV. As a consequence, the listed mass deficits are not a measure for the stability or binding energy of isolated nuclei, but for the whole atoms. This has very practical reasons, because it is very hard to totally ionize heavy elements, i.e. strip them of all of their electrons.

This practice is useful for other reasons, too: Stripping all the electrons from a heavy unstable nucleus (thus producing a bare nucleus) changes the lifetime of the nucleus, indicating that the nucleus cannot be treated independently (Experiments at the heavy ion accelerator GSI). This is also evident from phenomena like electron capture. Theoretically, in orbital models of heavy atoms, the electron orbits partially inside the nucleus (it doesn't orbit in a strict sense, but has a non-vanishing probability of being located inside the nucleus).

A nuclear decay happens to the nucleus, meaning that properties ascribed to the nucleus change in the event. In the field of physics the concept of "mass deficit" as a measure for "binding energy" means "mass deficit of the neutral atom" (not just the nucleus) and is a measure for stability of the whole atom.

Nuclear binding energy curve

In the periodic table of elements, the series of light elements from hydrogen up to sodium is observed to exhibit generally increasing binding energy per nucleon as the atomic mass increases. This increase is generated by increasing forces per nucleon in the nucleus, as each additional nucleon is attracted by other nearby nucleons, and thus more tightly bound to the whole.

The region of increasing binding energy is followed by a region of relative stability (saturation) in the sequence from magnesium through xenon. In this region, the nucleus has become large enough that nuclear forces no longer completely extend efficiently across its width. Attractive nuclear forces in this region, as atomic mass increases, are nearly balanced by repellent electromagnetic forces between protons, as the atomic number increases.

Finally, in elements heavier than xenon, there is a decrease in binding energy per nucleon as atomic number increases. In this region of nuclear size, electromagnetic repulsive forces are beginning to overcome the strong nuclear force attraction.

At the peak of binding energy, nickel-62 is the most tightly bound nucleus (per nucleon), followed by iron-58 and iron-56.^[9] This is the approximate basic reason why iron and nickel are very common metals in planetary cores, since they are produced profusely as end products in supernovae and in the final stages of silicon burning in stars. However, it is not binding energy per defined nucleon (as defined above), which controls which exact nuclei are made, because within stars, neutrons are free to convert to protons to release even more energy, per generic nucleon, if the result is a stable nucleus with a larger fraction of protons. In fact, it has been argued that photodisintegration of ⁶²Ni to form ⁵⁶Fe may be energetically possible in an extremely hot star core, due to this beta decay conversion of neutrons to protons.^[10] The conclusion is that at the pressure and temperature conditions in the cores of large stars, energy is released by converting all matter into ⁵⁶Fe nuclei (ionized atoms). (However, at high temperatures not all matter will be in the lowest energy state.) This energetic maximum should also hold for ambient conditions, say T = 298 K and p = 1 atm, for neutral condensed matter consisting of ⁵⁶Fe atoms—however, in these conditions nuclei of atoms are inhibited from fusing into the most stable and low energy state of matter.

It is generally believed that iron-56 is more common than nickel isotopes in the universe for mechanistic reasons, because its unstable progenitor nickel-56 is copiously made by staged build-up of 14 helium nuclei inside supernovas, where it has no time to decay to iron before being released into the interstellar medium in a matter of a few minutes, as the supernova explodes. However, nickel-56 then decays to cobalt-56 within a few weeks, then this radioisotope finally decays to iron-56 with a half life of about 77.3 days. The radioactive decay-powered light curve of such a process has been observed to happen in type II supernovae, such as SN 1987A. In a star, there are no good ways to create nickel-62 by alpha-addition processes, or else there would presumably be more of this highly stable nuclide in the universe.

Measuring the binding energy

The fact that the maximum binding energy is found in medium-sized nuclei is a consequence of the trade-off in the effects of two opposing forces that have different range characteristics. The attractive nuclear force (strong nuclear force), which binds protons and neutrons equally to each other, has a limited range due to a rapid exponential decrease in this force with distance. However, the repelling electromagnetic force, which acts between protons to force nuclei apart, falls off with distance much more slowly (as the inverse square of distance). For nuclei larger than about four nucleons in diameter, the additional repelling force of additional protons more than offsets any binding energy that results between further added nucleons as a result of additional strong force interactions. Such nuclei become increasingly less tightly bound as their size increases, though most of them are still stable. Finally, nuclei containing more than 209 nucleons (larger than about 6 nucleons in diameter) are all too large to be stable, and are subject to spontaneous decay to smaller nuclei.

Nuclear fusion produces energy by combining the very lightest elements into more tightly bound elements (such as hydrogen into helium), and nuclear fission produces energy by splitting the heaviest elements (such as uranium and plutonium) into more tightly bound elements (such as barium and krypton). Both processes produce energy, because middle-sized nuclei are the most tightly bound of all.

As seen above in the example of deuterium, nuclear binding energies are large enough that they may be easily measured as fractional mass deficits, according to the equivalence of mass and energy. The atomic binding energy is simply the amount of energy (and mass) released, when a collection of free nucleons are joined together to form a nucleus.

Nuclear binding energy can be computed from the difference in mass of a nucleus, and the sum of the masses of the number of free neutrons and protons that make up the nucleus. Once this mass difference, called the mass defect or mass deficiency, is known, Einstein's mass-energy equivalence formula E = mc² can be used to compute the binding energy of any nucleus. Early nuclear physicists used to refer to computing this value as a "packing fraction" calculation.

For example, the atomic mass unit (1 u) is defined as 1/12 of the mass of a ¹²C atom—but the atomic mass of a ¹H atom (which is a proton plus electron) is 1.007825 u, so each nucleon in ¹²C has lost, on average, about 0.8% of its mass in the form of binding energy.

Semiempirical formula for nuclear binding energy

For a nucleus with A nucleons, including Z protons and N neutrons, a semi-empirical formula for the binding energy (BE) per nucleon is:

where the coefficients are given by: ; ; ; ; .
The first term is called the saturation contribution and ensures that the binding energy per nucleon is the same for all nuclei to a first approximation. The term is a surface tension effect and is proportional to the number of nucleons that are situated on the nuclear surface; it is largest for light nuclei. The term is the Coulomb electrostatic repulsion; this becomes more important as increases. The symmetry correction term takes into account the fact that in the absence of other effects the most stable arrangement has equal numbers of protons and neutrons; this is because the n-p interaction in a nucleus is stronger than either the n-n or p-p interaction. The pairing term is purely empirical; it is + for even-even nuclei and - for odd-odd nuclei.

A graphical representation of the semi-
empirical binding energy formula. The
binding energy per nucleon in MeV (highest
numbers in dark red, in excess of 8.5 MeV
per nucleon) is plotted for various nuclides
 as a function of Z, the atomic number
 (y-axis), vs. N, the number of neutrons
 (x-axis). The highest numbers are seen for
Z = 26 (iron).

Example values deduced from experimentally measured atom nuclide masses

The following table lists some binding energies and mass defect values.^[11] Notice also that we use 1 u = (931.494028 ± 0.000023) MeV. To calculate the binding energy we use the formula Z (m_p + m_e) + N m_n - m_nuclide where Z denotes the number of protons in the nuclides and N their number of neutrons. We take m_p = 938.2723 MeV, m_e = 0.5110 MeV and m_n = 939.5656 MeV. The letter A denotes the sum of Z and N (number of nucleons in the nuclide). If we assume the reference nucleon has the mass of a neutron (so that all "total" binding energies calculated are maximal) we could define the total binding energy as the difference from the mass of the nucleus, and the mass of a collection of A free neutrons. In other words, it would be (Z + N) m_n - m_nuclide. The "total binding energy per nucleon" would be this value divided by A.

Most strongly bound nuclides atoms

nuclide	Z	N	mass excess	total mass	total mass / A	total binding energy / A	mass defect	binding energy	binding energy / A
56Fe	26	30	-60.6054 MeV	55.934937 u	0.9988372 u	9.1538 MeV	0.528479 u	492.275 MeV	8.7906 MeV
58Fe	26	32	-62.1534 MeV	57.932276 u	0.9988496 u	9.1432 MeV	0.547471 u	509.966 MeV	8.7925 MeV
60Ni	28	32	-64.472 MeV	59.93079 u	0.9988464 u	9.1462 MeV	0.565612 u	526.864 MeV	8.7811 MeV
62Ni	28	34	-66.7461 MeV	61.928345 u	0.9988443 u	9.1481 MeV	0.585383 u	545.281 MeV	8.7948 MeV

⁵⁶Fe has the lowest nucleon-specific mass of the four nuclides listed in this table, but this does not imply it is the strongest bound atom per hadron, unless the choice of beginning hadrons is completely free. Iron releases the largest energy if any 56 nucleons are allowed to build a nuclide—changing one to another if necessary, The highest binding energy per hadron, with the hadrons starting as the same number of protons Z and total nucleons A as in the bound nucleus, is ⁶²Ni. Thus, the true absolute value of the total binding energy of a nucleus depends on what we are allowed to construct the nucleus out of. If all nuclei of mass number A were to be allowed to be constructed of A neutrons, then Fe-56 would release the most energy per nucleon, since it has a larger fraction of protons than Ni-62. However, if nucleons are required to be constructed of only the same number of protons and neutrons that they contain, then nickel-62 is the most tightly bound nucleus, per nucleon.

Some light nuclides resp. atoms

nuclide	Z	N	mass excess	total mass	total mass / A	total binding energy / A	mass defect	binding energy	binding energy / A
n	0	1	8.0716 MeV	1.008665 u	1.008665 u	0.0000 MeV	0 u	0 MeV	0 MeV
1H	1	0	7.2890 MeV	1.007825 u	1.007825 u	0.7826 MeV	0.0000000146 u	0.0000136 MeV	13.6 eV
2H	1	1	13.13572 MeV	2.014102 u	1.007051 u	1.50346 MeV	0.002388 u	2.22452 MeV	1.11226 MeV
3H	1	2	14.9498 MeV	3.016049 u	1.005350 u	3.08815 MeV	0.0091058 u	8.4820 MeV	2.8273 MeV
3He	2	1	14.9312 MeV	3.016029 u	1.005343 u	3.09433 MeV	0.0082857 u	7.7181 MeV	2.5727 MeV

In the table above it can be seen that the decay of a neutron, as well as the transformation of tritium into helium-3, releases energy; hence, it manifests a stronger bound new state when measured against the mass of an equal number of neutrons (and also a lighter state per number of total hadrons). Such reactions are not driven by changes in binding energies as calculated from previously fixed N and Z numbers of neutrons and protons, but rather in decreases in the total mass of the nuclide/per nucleon, with the reaction. (Note that the Binding Energy given above for hydrogen-1 is the atomic binding energy, not the nuclear binding energy which would be zero.)

Eta Carinae

From Wikipedia, the free encyclopedia

Eta Carinae

Hubble Space Telescope image showing the bipolar Homunculus Nebula which surrounds Eta Carinae.
Observation data Epoch J2000      Equinox J2000
Constellation	Carina
Right ascension	10h 45m 03.591s[1]
Declination	−59° 41′ 04.26″[1]
Apparent magnitude (V)	−0.8 to 7.9[2]
Characteristics
Spectral type	F:I_pec_e[3] / O[4][5]
U−B color index	-0.45
B−V color index	0.61
Variable type	LBV[2] & binary
Astrometry
Radial velocity (Rv)	−25.0[6] km/s
Proper motion (μ)	RA: −7.6[1] mas/yr Dec.: 1.0[1] mas/yr
Absolute magnitude (MV)	-7 (current)
Details
Mass	120 / 30[7] M☉
Radius	~240[8] / 24[4] R☉
Luminosity	5,000,000 / <1 class="reference" id="cite_ref-verner_4-2" sup="">[4]

^[5] L_☉
Temperature
9,400^[9] / 37,200^[4] K
Age
<3 10="" sup="">6
^[5] years Orbit
Primary
Eta Carinae A
Companion
Eta Carinae B
Period (P)
2,022.7 days^[10]
(5.54 yr)
Semi-major axis (a)
15.4^[11] AU
Eccentricity (e)
0.9^[12]
Inclination (i)
~45^[11]°
Other designations

Foramen, Tseen She, 231G Carinae,^[13] HR 4210, CD−59°2620, HD 93308, SAO 238429, WDS 10451-5941, IRAS 10431-5925, GC 14799, CCDM J10451-5941

Database references
SIMBAD
data


Eta Carinae (η Carinae or η Car) is a stellar system containing at least two stars, about 7,500 light-years from the Sun in the direction of the constellation Carina. It is a member of the Trumpler 16 open cluster within the much larger Carina Nebula and currently has a combined bolometric luminosity of over five million times the Sun's.^[9]

Eta Carinae is circumpolar south of latitude 30°S, but is never visible north of latitude 30°N. It was first recorded as a 4th magnitude star, became the second brightest star in the sky before fading well below naked eye visibility, then brightened again and is now back at 4th magnitude.

The two main stars of the Eta Carinae system revolve in an eccentric orbit every 5.54 years. The primary is a peculiar star similar to a luminous blue variable (LBV) that initially had around 150 solar masses and has lost at least 30. Because of its mass and the stage of its life, it is expected to explode as a supernova or hypernova in the astronomically near future. This is currently the only star known to emit natural LASER light in ultraviolet wavelengths.^[14] The secondary is a hot star, probably class O, of approximately 30 solar masses and is itself a highly luminous star. The system is heavily obscured by the Homunculus Nebula, material ejected from the primary during its Great Eruption in the 19th century.

Observations

Eta Carinae is currently a 4th magnitude star, comfortably visible to the naked eye in dark skies.

Discovery

The earliest reliable record of Eta Carinae was made by Edmond Halley in 1677 when he recorded the star simply as "Sequens" (i.e. "following" relative to another star) within a new constellation Robur Carolinum. His Catalogus Stellarum Australium was published in 1679.^[15]

There are some possible earlier observations of Eta Carinae. There are speculative reports from antiquity that may relate to Eta Carinae, but no reliable observations. Most observations of bright stars in the southern constellations in the 16th century fail to record Eta Carinae. Pieter Keyser described a fourth magnitude star at approximately the correct position in the late 16th century, but Frederick de Houtman's catalogue from 1603 does not include Eta Carinae among the other fourth magnitude stars in the region.^[15]

In 1751 Nicolas Louis de Lacaille mapped the stars of Argo Navis and Robur Carolinum into separate smaller constellations and gave the brighter members Greek alphabet Bayer designations. Eta was placed within the keel portion of the ship named as the new constellation Carina.^[16]

Surroundings

The Carina Nebula. Eta Carinae is the brightest star, just left of centre.

Annotated image of Carina Nebula

Eta Carinae lies within the scattered stars of the Trumpler 16 open cluster. All the other members are well below naked eye visibility, although WR 25 is another extremely massive luminous star. Trumpler 16 and its neighbour Trumpler 14 are the main star clusters of the Carina OB1 association, one of the two main stellar associations of the Carina Nebula, together with Carina OB2.

Eta Carinae is enclosed by the Homunculus Nebula, a reflection nebula lit mainly by Eta Carinae itself.^[17] The Homunculus Nebula is composed mainly of dust which condensed from the debris ejected during the Great Eruption event in the mid nineteenth century. The nebula consists primarily of two polar lobes aligned with the rotation axis of the star, plus an equatorial "skirt". Closer studies show many fine details: a Little Homunculus within the main nebula, probably formed by the 1890 eurption; a jet; fine streams and knots of material, especially noticeable in the skirt region; and three Weigelt Blobs, dense gas condensations very close to the star itself.^[14][18]

Brightness

Combined UV and visual light image of Eta Carinae

Halley gave an approximate apparent magnitude of "4" at the time of discovery, which has been calculated as magnitude 3.3 on the modern scale. The handful of possible earlier sightings suggest that Eta Carinae was not significantly brighter than this for much of the 17th century.^[15]

There are further sporadic observations over the next 70 years showing that Eta Carinae was probably around or below 4th magnitude, until Lacaille's recorded it at second magnitude in 1751.^[16] It is unclear whether Eta Carinae varied significantly in brightness over the next 50 years, with occasional observations such as William Burchell at 4th magnitude in 1815, but it is uncertain whether these are just re-recordings of earlier observations. In 1827 Burchell specifically noted its unusual brightness at 1st magnitude.^[15]

In the 1830s John Herschel made a detailed series of accurate measurements showing Eta Carinae consistently around magnitude 1.2. However at the end of 1837, it brightened suddenly to magnitude 0 before dropping slightly to magnitude 0.6.^[19]

In summary, the brightness of Eta Carinae increased from around 4th magnitude to 1st magnitude over about 150 years, possibly erratically, before brightening dramatically. This was the start of the Great Eruption, which by 1843 saw Eta Carinae become the second brightest star in the sky after Sirius. To put the relationship in perspective, Sirus is nearly a thousand times closer, but only appears 40% brighter than Eta Carinae at its peak. Particular peaks in 1827, 1838, and 1843 may have been related to the periastron passage of the binary orbit.^[20] From 1845 to 1856, the brightness decreased by around 0.1 magnitudes per year, but with possible rapid and large fluctuations.^[15]

From 1857 the brightness decreased rapidly until it faded below naked eye visibility by 1886. This has been calculated to be due to the condensation of dust in the ejected material surrounding the star rather than an intrinsic change in luminosity.^[21] There was a brightening from 1887 - 1895, peaking at about magnitude 6.2 then dimming rapidly to about magnitude 7.5. This appeared to be a smaller copy of the Great Eruption, expelling material that formed the Little Homunculus and Weigelt Blobs.^[22][23]

For the first half of the 20th century, Eta Carinae appeared to have settled at a constant brightness at 8th magnitude, but in 1953 it was noted to have brightened again to magnitude 6.5.^[24] The brightening continued steadily, but with fairly regular variations of a few tenths of a magnitude that were later identified as having a 5.54 year period.^[20] A sudden doubling of brightness was observed in 1998–1999 bringing it back to naked eye visibility. As of 2012, the visual magnitude was 4.6.^[25]

The brightness doesn't always vary consistently at different wavelengths, and does not always consistently follow the 5.5 year cycle.^[5][26]

Spectrum

Hubble composite of Eta Carinae showing the unusual emission spectrum taken using STIS

The spectrum of Eta Carinae is peculiar and variable. The earliest observations of the spectrum in 1893 are described as similar to an F5 star, but with a few emission lines. Analysis to modern spectral standards suggests an early F spectral type.^[27] By 1895, the spectrum consisted mostly of strong emission lines, with the absorption lines present but largely obscured by emission. The lines vary greatly in width and profile.^[28][29]

Direct spectral observations do not begin until after the Great Eruption, but light echoes from the eruption were detected using the U.S. National Optical Astronomy Observatory's Blanco 4-meter telescope at the Cerro Tololo Inter-American Observatory. Analysis of the reflected spectra indicated the light was emitted when Eta Carinae had the appearance of a 5,000 K G2-to-G5 supergiant, some 2,000 K cooler than expected from other supernova impostor events.^[30] Further light echo observations show that following the peak brightness of the Great Eruption the spectrum developed prominent P Cygni profiles and CN molecular bands. These indicate that the star, or the expanding cloud of ejected material, had cooled further and may have been colliding with circumstellar material in a similar way to a type IIn supernova.^[31]

In the second half of the 20th century, infra-red and ultra-violet spectra became available, as well as much higher resolution visual spectra. The spectrum continued to show complex and baffling features, with much of the energy from the central star being recycled into the infra-red by surrounding dust, some reflection of light from the star from dense localised objects in the circumstellar material, but with obvious high ionisation features indicative of very high temperatures. The line profiles are complex and variable, indicating a number of absorption and emission features at various velocities relative to the central star.^[32][33]

The 5.5 year orbital cycle produces strong spectral changes at periastron that are known as "spectroscopic event"s. Certain wavelengths of radiation suffer eclipses, either due to actual occultation by one of the stars, or due to passage within opaque portions of the complex stellar winds. Despite being ascribed to orbital rotation, these events vary significantly from cycle to cycle. These changes have become stronger since 2003 and it is generally believed that longterm secular changes in the stellar winds or previously ejected material may be the culmination of a return to the state of the star prior to its Great Eruption.^[26][34][35]

High energy radiation

X-Rays around Eta Carinae (red is low energy, blue higher)

Several x-ray and gamma-ray sources have been detected around Eta Carinae, for example 4U 1037–60 in the 4th Uhuru catalogue and 1044–59 in the HEAO-2 catalog. The earliest detection of x-rays in the Eta Carinae region was from the Terrier-Sandhawk rocket, ^[36] followed by Ariel 5,^[37] OSO 8,^[38] and Uhuru^[39] sightings.

More detailed observations were made with Einstein,^[40] ROSAT,^[41] ASCA,^[42] and Chandra. There are multiple sources at various wavelengths right across the high energy electromagnetic spectrum: hard x-rays and gamma rays within 1 light month of the Eta Carinae; hard x-rays from a central region about 3 light months wide; a distinct partial ring "horse-shoe" structure in low energy x-rays 0.67 pc across corresponding to the main shockfront from the Great Eruption; diffuse x-ray emission across the whole area of the Homunculus; and numerous condensations and arcs outside the main ring.^{[43][43][44][45][46]}

All the high energy emission associated with Eta Carinae varies during the orbital cycle. A spectroscopic minimum, or X-ray eclipse, occurred in July and August 2003 and similar events in 2009 and 2014 have been intensively observed.^[47] The highest energy gamma-rays above 100MeV detected by AGILE show strong variability, while lower energy gamma-rays observed by Fermi show little variability.^[43][48]

Radio emission

Eta Carinae has been detected at various radio wavelengths including EHF (millimeter wave), SHF, and UHF. The detected radiation appears to be both thermal emission from warm gas and free-free emission from ionised gas. The free-free emission in particular varies during the orbital cycle, with significant dips during the periastron passage Spectroscopic events.^[49][50]

Properties

X-ray, optical and infrared images of Eta Carinae (August 26, 2014).

The Eta Carinae stellar system is currently one of the most massive that can be studied in great detail. Until recently Eta Carinae was thought to be the most massive single star, but in 2005 the system's binary nature was confirmed.^[51] Unfortunately, both component stars are largely obscured by circumstellar material ejected from Eta Carinae A and basic properties such as their temperatures and luminosities can only be inferred. Rapid changes to the stellar wind in the 21st century suggest that the star itself may be revealed as dust from the great eruption finally clears.^[52]

Classification

Eta Carinae A is classified as a luminous blue variable (LBV) due to peculiarities in its pattern of brightening and dimming. This type of variable star is characterised by irregular changes from a high temperature quiescent state to a low temperature eruptive state at roughly constant luminosity. LBVs in the quiescent state lie on a narrow S Doradus instability strip, with more luminous stars being hotter. In eruption all LBVs have about the same temperature near 8,000K. LBVs in eruption are visually brighter than when quiescent although the bolometric luminosity is unchanged.

Eta Carinae A is not a typical LBV. It is more luminous than any other LBV in the Milky Way although possibly comparable to other Supernova Imposters detected in external galaxies. It doesn't currently lie on the S Doradus instability strip, although it is unclear what the temperature or spectral type of the underlying star actually is. The 1890 eruption may have been fairly typical of LBV eruptions, with an early F spectral type, and it has been estimated that the star may currently have an opaque stellar wind forming a pseudo-photosphere with a temperature of 9,000K - 14,000K which would be typical for an LBV in eruption.^[21]

The Great Eruption event of Eta Carinae A has only been observed in a handful of other possible LBVs in external galaxies. It isn't clear if this is something that only a very few of the most massive LBVs undergo, something that is caused by a close companion star, or a very brief but common phase for massive stars.

Eta Carinae B is a massive luminous hot star, but little else is known. From certain high excitation spectral lines that ought not to be produced by the primary, it is thought that Eta Carinae B is a young O-type star. Most authors suggest it is a somewhat evolved star such as a supergiant or giant, although a Wolf-Rayet star cannot be ruled out.^[51]

Mass

The masses of stars are difficult to measure except by determination of a binary orbit. Eta Carinae is a binary system, but certain key information about the orbit is not known accurately. Several models of the system use masses of 120-160 M_☉ and 30-60 M_☉ for the primary and secondary respectively. Eta Carinae A has clearly lost a great deal of mass since it formed and was initially 150-180 M_☉.^[7]

Mass loss

HST image of the Homunculus Nebula, inset is a VLT NACO infra-red image of Eta Carinae

Mass loss is one of the most intensively studied aspects of massive star research. Put simply, using observed mass loss rates in the best models of stellar evolution do not reproduce the observed distribution of evolved massive stars such as Wolf-Rayets, the number and types of core collapse supernovae, or their progenitors. To match those observations, the models require much higher mass loss rates. Eta Carinae A has one of the highest known mass loss rates, currently around 10^-3 M_☉/year, and is an obvious candidate for study.^[7]

Eta Carinae A is losing so much mass due to its extreme luminosity and relatively low surface gravity that the stellar wind is entirely opaque and appears as a pseudo-photosphere. This optically dense surface hides the true physical surface of the star. During the Great Eruption the mass loss rate was a thousand times higher, around 1 M_☉/year sustained for ten years or more. The total mass loss during the eruption was 10-20 M_☉ with much of it now forming the Homunculus Nebula. The smaller 1890 eruption produced the Little Homunculus Nebula, much smaller and only about 0.1 M_☉.^[8] The bulk of the mass loss occurs in a wind with a terminal velocity of about 400 km/s, but some material is seen at higher velocities, up to 3,200 km/s, possibly material blown from the accretion disk by the secondary star.^[53]

Eta Carinae B is presumably also losing mass via a thin fast stellar wind, but this cannot be detected directly. Models of the radiation observed from interactions between the winds of the two stars show a mass loss rate of the order of 10^-6 M_☉/year, typical of a hot O class star. For a portion of the highly eccentric orbit, it actually gains material from the primary via an accretion disk. During the Great Eruption of the primary, the secondary accreted several M_☉, producing strong jets which formed the bipolar shape of the Homunculus Nebula.^[7]

Luminosity

The stars of the Eta Carinae system are completely obscured by dust and opaque stellar winds. The total electromagnetic radiation across all wavelengths for both stars combined is several million L_☉. The best estimate for the luminosity of the primary is 5 million L_☉. The luminosity of Eta Carinae B is particularly uncertain, probably several hundred thousand L_☉ and almost certainly no more than 1 million L_☉. Due to the surrounding dust, 90% of the radiation from the stars reaches us as infra-red and Eta Carinae is the brightest IR source outside the solar system.^[21]

The most notable feature of Eta Carinae is its giant eruption or supernova impostor event, which originated in the primary star and was observed around 1843. In a few years, it produced almost as much visible light as a faint supernova explosion, but the star survived. It is estimated that at peak brightness it was around 25 million times more luminous than the sun.^[30] Other supernova impostors have been seen in other galaxies, for example the possible false supernova SN 1961v in NGC 1058^[54] and SN 2006jc in UGC 4904,^[55] which produced a false supernova noted in October 2004. Significantly SN 2006jc was destroyed in a supernova explosion two years later on October 9, 2006.

Temperature

Eta Carinae taken with the NACO near-IR adaptive optics instrument of ESO's Very Large Telescope.

The temperature of Eta Carinae B can be estimated with some accuracy due to spectral lines that are only likely to be produced by a star around 37,000 K.^[5]

The temperature of the primary star is more uncertain. For many years it was expected to be over 30,000 K due to the presence of the high temperature spectral lines now attributed to the secondary star, although this conflicted with other spectral characteristics that ought only to be found in cooler stars. That conflict is now resolved, and Eta Carinae A, or at least what we can see of it, is accepted to be considerably cooler than Eta Carinae B. The star is likely to have been an early B hypergiant with a temperature of 20,000 K - 25,000 K at the time of its dscovery by Halley. An effective temperature based on its luminosity today would also be around 20,000 K - 25,000 K, but the hints of light directly from the star itself via some dense nebular features suggest a much cooler star at 9,000 K - 14,000 K. This cooler temperature may be a pseudo-photosphere formed where the opaque stellar wind starts to become transparent. Very recent observations show dramatic changes in the stellar wind and a possible increase in the temperature of any pseudo-photosphere, but it is still largely shrouded in dust shifting most of the light output into the infra-red.^[25][34][52]

The powerful stellar winds from the two stars collide and produce temperatures as high as 7001600000000000000♠60 MK, which is the source of the hard x-rays and gamma-rays close to the stars. Further out, expanding gases from the Great Eruption collide with interstellar material and are heated to around 7001600000000000000♠60 MK, producing less energetic x-rays seen in a ring shape.

Size

The size of the two main stars in the Eta Carinae system is difficult to determine precisely because neither star can be seen directly. Eta Carinae B is likely to have a well-defined photosphere and its radius can be estimated from the assumed type of star. An O supergiant of 933,000 L_☉ with a temperature of 37,200K has an effective radius of 23.6 R_☉.^[4]

The size of Eta Carinae A is not even well defined. It has an optically dense stellar wind so the typical definition of a star's surface being approximately where it becomes opaque gives a very different result to where a more traditional definition of a surface might be. One study calculated a radius of 60 R_☉ for a hot "core" of 35,000K at optical depth 150, near the sonic point or very approximately what might be called a physical surface, but over 800 R_☉ for optical depth 0.67, the visible surface of the stellar wind.^[56]

Rotation

Rotation rates of massive stars have a critical influence on their evolution and eventual death. The rotation rate of the Eta Carinae stars cannot be measured directly because their surfaces cannot be seen. Single massive stars spin down quickly due to braking from their string winds, but there are hints that both Eta Carinae A and B are fast rotators. One or both could have been spun up by binary interaction, for example accretion onto the secondary, and orbital dragging on the primary.^[57]

Future prospects

With their disproportionately high luminosities, very large stars such as Eta Carinae use up their fuel very quickly. Eta Carinae is expected to explode as a supernova or hypernova some time within the next million years or so. As its current age and evolutionary path are uncertain, however, it could explode within the next several millennia or even in the next few years. LBVs such as Eta Carinae may be a stage in the evolution of the most massive stars; the prevailing theory now holds that they will exhibit extreme mass loss and become Wolf-Rayet stars before they go supernova, if they are unable to hold their mass to explode as a hypernova.^[58]

More recently, another possible Eta Carinae analogue was observed: SN 2006jc, some 77 million light years away in UGC 4904, in the constellation of Lynx.^[55] Its brightened appearance was noted on 20 October 2004, and was reported by amateur astronomer Koichi Itagaki as a supernova. However, although it had indeed exploded, hurling 0.01 solar masses (~20 Jupiters) of material into space, it had survived, before finally exploding nearly two years later as a Mag 13.8 type Ib supernova, seen on 9 October 2006. Its earlier brightening was a supernova impostor event.

The similarity between Eta Carinae and SN 2006jc has led Stefan Immler of NASA's Goddard Space Flight Center to suggest that Eta Carinae could explode in our lifetime, or even in the next few years.^[55] However, Stanford Woosley of the University of California in Santa Cruz disagrees with Immler’s suggestion, and says it is likely that Eta Carinae is at an earlier stage of evolution, and that there are still several stages of nuclear burning to go before the star runs out of fuel. When it does occur, the supernova will be brighter than Venus but not as bright as a full moon.^[59]

In NGC 1260, a spiral galaxy in the constellation of Perseus some 238 million light years from earth, another analogue star explosion, supernova SN 2006gy, was observed on September 18, 2006. A number of astronomers modelling supernova events have suggested that the explosion mechanism for SN 2006gy may be very similar to the fate that awaits Eta Carinae.^{[citation needed]}

Possible effects on Earth

One theory of Eta Carinae's ultimate fate is collapsing to form a black hole - energy released as jets along the axis of rotation forms gamma ray bursts.

It is possible that the Eta Carinae hypernova or supernova, when it occurs, could affect Earth, which is about 7,500 light years from the star. It is unlikely, however, to affect terrestrial lifeforms directly, as they will be protected from gamma rays by the atmosphere, and from some other cosmic rays by the magnetosphere. The damage would likely be restricted to the upper atmosphere, the ozone layer, spacecraft, including satellites, and any astronauts in space. However, at least one paper has projected that complete loss of the Earth's ozone layer is a plausible consequence of a nearby supernova, which would result in a significant increase in surface UV radiation reaching the Earth's surface from our own Sun.^[60] A supernova or hypernova produced by Eta Carinae would probably eject a gamma ray burst (GRB) out from both polar areas of its rotational axis. Calculations show that the deposited energy of such a GRB striking the Earth's atmosphere would be equivalent to one kiloton of TNT per square kilometer over the entire hemisphere facing the star, with ionizing radiation depositing ten times the lethal whole body dose to the surface.^[61] This catastrophic burst would probably not hit Earth, though, because the rotation axis does not currently point towards our solar system. If Eta Carinae is a binary system, this may affect the future intensity and orientation of the supernova explosion that it produces, depending on the circumstances.^[51]

Cultural significance

In traditional Chinese astronomy, Eta Carinae has the names Tseen She (from the Chinese 天社 [Mandarin: tiānshè] "Heaven's altar") and Foramen. It is also known as 海山二 (Hǎi Shān èr, English: the Second Star of Sea and Mountain),^[62] referring to Sea and Mountain, an asterism that Eta Carinae forms with s Carinae, λ Centauri and λ Muscae.^[63]

In 2010, astronomers Duane Hamacher and David Frew from Macquarie University in Sydney showed that the Boorong Aboriginal people of northwestern Victoria, Australia, witnessed the outburst of Eta Carinae in the 1840s and incorporated it into their oral traditions as Collowgulloric War, the wife of War (Canopus, the Crow – wɑː).^[64] This is the only definitive indigenous record of Eta Carinae's outburst identified in the literature to date.

Computers that mimic the function of the brain

Apr 06, 2015 by Amanda Morri

Original link:  http://phys.org/news/2015-04-mimic-function-brain.html?utm_source=menu&utm_medium=link&utm_campaign=item-menu

GB migration. Credit: Nature Nanotechnology (2015) doi:10.1038/nnano.2015.56

Researchers are always searching for improved technologies, but the most efficient computer possible already exists. It can learn and adapt without needing to be programmed or updated. It has nearly limitless memory, is difficult to crash, and works at extremely fast speeds. It's not a Mac or a PC; it's the human brain. And scientists around the world want to mimic its abilities.

Both academic and industrial laboratories are working to develop computers that operate more like the human brain. Instead of operating like a conventional, digital system, these new devices could potentially function more like a network of neurons.

"Computers are very impressive in many ways, but they're not equal to the mind," said Mark Hersam, the Bette and Neison Harris Chair in Teaching Excellence in Northwestern University's McCormick School of Engineering. "Neurons can achieve very complicated computation with very low power consumption compared to a digital computer."

A team of Northwestern researchers, including Hersam, has accomplished a new step forward in electronics that could bring brain-like computing closer to reality. The team's work advances memory resistors, or "memristors," which are resistors in a circuit that "remember" how much current has flowed through them.

The research is described in the April 6 issue of Nature Nanotechnology. Tobin Marks, the Vladimir N. Ipatieff Professor of Catalytic Chemistry, and Lincoln Lauhon, professor of materials science and engineering, are also authors on the paper. Vinod Sangwan, a postdoctoral fellow co-advised by Hersam, Marks, and Lauhon, served as first author. The remaining co-authors—Deep Jariwala, In Soo Kim, and Kan-Sheng Chen—are members of the Hersam, Marks, and/or Lauhon research groups.

"Memristors could be used as a memory element in an integrated circuit or computer," Hersam said. "Unlike other memories that exist today in modern electronics, memristors are stable and remember their state even if you lose power."

Current computers use random access memory (RAM), which moves very quickly as a user works but does not retain unsaved data if power is lost. Flash drives, on the other hand, store information when they are not powered but work much slower. Memristors could provide a memory that is the best of both worlds: fast and reliable. But there's a problem: memristors are two-terminal electronic devices, which can only control one voltage channel. Hersam wanted to transform it into a three-terminal device, allowing it to be used in more complex electronic circuits and systems.

Hersam and his team met this challenge by using single-layer molybdenum disulfide (MoS2), an atomically thin, two-dimensional nanomaterial semiconductor. Much like the way fibers are arranged in wood, atoms are arranged in a certain direction—called "grains"—within a material. The sheet of MoS2 that Hersam used has a well-defined grain boundary, which is the interface where two different grains come together.

"Because the atoms are not in the same orientation, there are unsatisfied chemical bonds at that interface," Hersam explained. "These grain boundaries influence the flow of current, so they can serve as a means of tuning resistance."

When a large electric field is applied, the grain boundary literally moves, causing a change in resistance. By using MoS2 with this grain boundary defect instead of the typical metal-oxide-metal memristor structure, the team presented a novel three-terminal memristive device that is widely tunable with a gate electrode.

"With a memristor that can be tuned with a third electrode, we have the possibility to realize a function you could not previously achieve," Hersam said. "A three-terminal memristor has been proposed as a means of realizing brain-like computing. We are now actively exploring this possibility in the laboratory."

Explore further: New process isolates promising material molybdenum disulfide

More information: Gate-tunable memristive phenomena mediated by grain boundaries in single-layer MoS2, Nature Nanotechnology (2015) DOI: 10.1038/nnano.2015.56

Journal reference: Nature Nanotechnology
Provided by Northwestern University