p-process

The term p-process (p is for proton) is used in two ways in the scientific literature concerning the astrophysical origin of the elements (nucleosynthesis). Originally it referred to a proton capture process which is the source of certain, naturally occurring, proton-rich isotopes of the elements from selenium to mercury.[1][2] These nuclides are called p-nuclei and their origin is still not completely understood. Although it was shown that the originally suggested process cannot produce the p-nuclei, later on the term p-process was sometimes used to generally refer to any nucleosynthesis process supposed to be responsible for the p-nuclei.[3]

Often, the two meanings are confused. Recent scientific literature therefore suggests to use the term p-process only for the actual proton capture process, as it is customary with other nucleosynthesis processes in astrophysics.[4]

The proton capture p-process

Proton-rich nuclides can be produced by sequentially adding one or more protons to an atomic nucleus. Such a nuclear reaction of type (p,γ) is called proton capture reaction. By adding a proton to a nucleus, the element is changed because the chemical element is defined by the proton number of a nucleus. At the same time the ratio of protons to neutrons is changed, resulting in a proton-richer isotope of the next element. This led to the original idea for the production of p-nuclei: free protons (the nuclei of hydrogen atoms are present in stellar plasmas) should be captured on heavy nuclei (seed nuclei) also already present in the stellar plasma (previously produced in the s- and/or r-process).[1][2]

Such proton captures on stable nuclides (or nearly stable), however, are not very efficient in producing p-nuclei, especially the heavier ones, because the electric charge increases with each added proton, leading to an increased repulsion of the next proton to be added, according to Coulomb's law. In the context of nuclear reactions this is called a Coulomb barrier. The higher the Coulomb barrier the more kinetic energy a proton requires to get close to a nucleus and be captured by it. The average energy of the available protons is given by the temperature of the stellar plasma. Even if this temperature could be increased arbitrarily (which is not the case in stellar environments), protons would be removed faster from a nucleus by photodisintegration than they could be captured at high temperature. A possible alternative would be to have a very large number of protons available to increase the effective number of proton captures per second without having to raise the temperature too much. Such conditions, however, are not found in core-collapse supernovae which were supposed to be the site of the p-process.[3][4]

Proton captures at extremely high proton densities are called rapid proton capture processes. They are distinct from the p-process not only by the required high proton density but also by the fact that very short-lived radionuclides are involved and the reaction path is located close to the proton drip line. Rapid proton capture processes are the rp-process, the νp-process, and the pn-process.

History

The term p-process was originally proposed in the famous "B2FH paper" in 1957. The authors assumed that this process was solely responsible for the p-nuclei and proposed that it occurs in the hydrogen-shell (see also stellar evolution) of a star exploding as a type II supernova.[1] It was shown later that the required conditions are not found in such supernovae.[5]

At the same time as B2FH, Alastair Cameron independently realized the necessity to add another nucleosynthesis process to neutron capture nucleosynthesis but simply mentioned proton captures without assigning a special name to the process. He also thought about alternatives, for example photodisintegration (called the γ-process today) or a combination of p-process and photodisintegration.[2]

See also

References

  1. ^ a b c Burbidge, E. M.; Burbidge, G. R.; Fowler, W. A.; Hoyle, F. (1957). "Synthesis of the Elements in Stars". Reviews of Modern Physics. 29 (4): 547–650. Bibcode:1957RvMP...29..547B. doi:10.1103/RevModPhys.29.547.
  2. ^ a b c Cameron, A. G. W. (1957). "Nuclear Reactions in Stars and Nucleogenesis". Publications of the Astronomical Society of the Pacific. 69 (408): 201–222. Bibcode:1957PASP...69..201C. doi:10.1086/127051. JSTOR 40676435.
  3. ^ a b Arnould, M.; Goriely, S. (2003). "The p-Process of Stellar Nucleosynthesis: Astrophysics and Nuclear Physics Status". Physics Reports. 384 (1–2): 1–84. Bibcode:2003PhR...384....1A. doi:10.1016/S0370-1573(03)00242-4.
  4. ^ a b Rauscher, T. (2010). "Origin of p-Nuclei in Explosive Nucleosynthesis". Proceedings of Science. NIC XI (59). arXiv:1012.2213. Bibcode:2010arXiv1012.2213R.
  5. ^ Audouze, J.; Truran, J. W. (1975). "p-Process Nucleosynthesis in Postshock Supernova Envelope Environments" (PDF). The Astrophysical Journal. 202 (1): 204–213. Bibcode:1975ApJ...202..204A. doi:10.1086/153965.
Autoregressive model

In statistics, econometrics and signal processing, an autoregressive (AR) model is a representation of a type of random process; as such, it is used to describe certain time-varying processes in nature, economics, etc. The autoregressive model specifies that the output variable depends linearly on its own previous values and on a stochastic term (an imperfectly predictable term); thus the model is in the form of a stochastic difference equation. Together with the moving-average (MA) model, it is a special case and key component of the more general ARMA and ARIMA models of time series, which have a more complicated stochastic structure; it is also a special case of the vector autoregressive model (VAR), which consists of a system of more than one interlocking stochastic difference equation in more than one evolving random variable.

Contrary to the moving-average model, the autoregressive model is not always stationary as it may contain a unit root.

Box–Jenkins method

In time series analysis, the Box–Jenkins method, named after the statisticians George Box and Gwilym Jenkins, applies autoregressive moving average (ARMA) or autoregressive integrated moving average (ARIMA) models to find the best fit of a time-series model to past values of a time series.

Dime (Canadian coin)

In Canada, a dime is a coin worth ten cents. It has been the physically smallest Canadian coin since 1922, smaller even than the penny despite its higher face value. According to the Royal Canadian Mint, the official national term of the coin is the 10-cent piece, but in practice, the term dime predominates in English-speaking Canada. It is nearly identical in size to the American dime, but unlike its counterpart, the Canadian dime is magnetic due to a distinct metal composition: from 1968 to 1999 it was composed entirely of nickel, and since 2000 it has had a high steel content.

Currently the dime has, as with all Canadian coins, a portrait of Queen Elizabeth II on the obverse. The reverse contains a representation of the Bluenose, a famous Canadian schooner. The artist, Emanuel Hahn, used three ships including the Bluenose as his models, so the ship design is actually a composite. The coin is produced by the Royal Canadian Mint at its facility in Winnipeg.

The word dime comes from the French word dîme, meaning "tithe" or "tenth part", from the Latin decima [pars].

IEC 61355

The standard IEC 61355-1 Classification and designation of documents for plants, systems and equipment describes rules and guidelines for the uniform classification and identification of documents based on their characteristic content of information.

It is applied for all documents within the life cycle of a technical products like plants, systems or equipment. It also includes non-technical documents. The main application is the construction, erection and operation of chemical plants and power plants, where the number of documents may sum up to some 100,000 documents.

Neutron capture

Neutron capture is a nuclear reaction in which an atomic nucleus and one or more neutrons collide and merge to form a heavier nucleus. Since neutrons have no electric charge, they can enter a nucleus more easily than positively charged protons, which are repelled electrostatically.Neutron capture plays an important role in the cosmic nucleosynthesis of heavy elements. In stars it can proceed in two ways: as a rapid (r-process) or a slow process (s-process). Nuclei of masses greater than 56 cannot be formed by thermonuclear reactions (i.e. by nuclear fusion), but can be formed by neutron capture.

Neutron capture on protons yields a line at 2.223 MeV predicted and commonly observed in solar flares.

Nuclear reaction

In nuclear physics and nuclear chemistry, a nuclear reaction is semantically considered to be the process in which two nuclei, or else a nucleus of an atom and a subatomic particle (such as a proton, neutron, or high energy electron) from outside the atom, collide to produce one or more nuclides that are different from the nuclide(s) that began the process. Thus, a nuclear reaction must cause a transformation of at least one nuclide to another. If a nucleus interacts with another nucleus or particle and they then separate without changing the nature of any nuclide, the process is simply referred to as a type of nuclear scattering, rather than a nuclear reaction.

In principle, a reaction can involve more than two particles colliding, but because the probability of three or more nuclei to meet at the same time at the same place is much less than for two nuclei, such an event is exceptionally rare (see triple alpha process for an example very close to a three-body nuclear reaction). "Nuclear reaction" is a term implying an induced changing in a nuclide, and thus it does not apply to any type of radioactive decay (which by definition is a spontaneous process).Natural nuclear reactions occur in the interaction between cosmic rays and matter, and nuclear reactions can be employed artificially to obtain nuclear energy, at an adjustable rate, on demand. Perhaps the most notable nuclear reactions are the nuclear chain reactions in fissionable materials that produce induced nuclear fission, and the various nuclear fusion reactions of light elements that power the energy production of the Sun and stars.

Oppenheimer–Phillips process

The Oppenheimer–Phillips process or strip reaction is a type of deuteron-induced nuclear reaction. In this process the neutron half of an energetic deuteron (a stable isotope of hydrogen with one proton and one neutron) fuses with a target nucleus, transmuting the target to a heavier isotope while ejecting a proton. An example is the nuclear transmutation of carbon-12 to carbon-13.

The process allows a nuclear interaction to take place at lower energies than would be expected from a simple calculation of the Coulomb barrier between a deuteron and a target nucleus. This is because, as the deuteron approaches the positively charged target nucleus, it experiences a charge polarization where the "proton-end" faces away from the target and the "neutron-end" faces towards the target. The fusion proceeds when the binding energy of the neutron and the target nucleus exceeds the binding energy of the deuteron and a proton is then repelled from the new, heavier, nucleus.

P-nuclei

p-nuclei (p stands for proton-rich) are certain proton-rich, naturally occurring isotopes of some elements between selenium and mercury inclusive which cannot be produced in either the s- or the r-process.

Partial autocorrelation function

In time series analysis, the partial autocorrelation function (PACF) gives the partial correlation of a stationary time series with its own lagged values, regressed the values of the time series at all shorter lags. It contrasts with the autocorrelation function, which does not control for other lags.

This function plays an important role in data analysis aimed at identifying the extent of the lag in an autoregressive model. The use of this function was introduced as part of the Box–Jenkins approach to time series modelling, whereby plotting the partial autocorrelative functions one could determine the appropriate lags p in an AR (p) model or in an extended ARIMA (p,d,q) model.

Pentium

Pentium is a brand used for a series of x86 architecture-compatible microprocessors produced by Intel since 1993. In their form as of November 2011, Pentium processors are considered entry-level products that Intel rates as "two stars", meaning that they are above the low-end Atom and Celeron series, but below the faster Core i3, i5, i7, i9, and workstation Xeon series.

As of 2017, Pentium processors have little more than their name in common with earlier Pentiums, which were Intel's flagship processor for over a decade until the introduction of the Intel Core line in 2006. They are based on both the architecture used in Atom and that of Core processors. In the case of Atom architectures, Pentiums are the highest performance implementations of the architecture. Pentium processors with Core architectures prior to 2017 were distinguished from the faster, higher-end i-series processors by lower clock rates and disabling some features, such as hyper-threading, virtualization and sometimes L3 cache.

The name Pentium is originally derived from the Greek word penta (πεντα), meaning "five", a reference to the prior numeric naming convention of Intel's 80x86 processors (8086–80486), with the Latin ending -ium.

In 2017, Intel split Pentium into two line-ups. Pentium Silver aiming for low-power devices and shares architecture with Atom and Celeron. Pentium Gold aiming for entry-level desktop and using existing architecture, such as Kaby Lake or Coffee Lake.

Photodisintegration

Photodisintegration (also called phototransmutation) is a nuclear process in which an atomic nucleus absorbs a high-energy gamma ray, enters an excited state, and immediately decays by emitting a subatomic particle. The incoming gamma ray effectively knocks one or more neutrons, protons, or an alpha particle out of the nucleus. The reactions are called (γ,n), (γ,p), and (γ,α).

Photodisintegration is endothermic (energy absorbing) for atomic nuclei lighter than iron and sometimes exothermic (energy releasing) for atomic nuclei heavier than iron. Photodisintegration is responsible for the nucleosynthesis of at least some heavy, proton-rich elements via the p-process in supernovae.

Proton capture

Proton capture is a nuclear reaction in which an atomic nucleus and one or more protons collide and merge to form a heavier nucleus.

Since protons have positive electric charge, they are repelled electrostatically by the positively charged nucleus. Therefore, it is more difficult for protons to enter the nucleus compared to neutrally charged neutrons .

Proton capture plays an important role in the cosmic nucleosynthesis of proton rich isotopes. In stars it can proceed in two ways: as a rapid (rp-process) or a slow process (p-process).

R-process

The rapid neutron-capture process, or so-called r-process, is a set of nuclear reactions that in nuclear astrophysics is responsible for the creation of approximately half of the atomic nuclei heavier than iron; the "heavy elements". The other half are produced by the p-process and s-process. The r-process usually synthesizes all of the two most neutron-rich stable isotopes of each heavy element.

The heavy elements typically have six to ten stable isotopes. Chemical elements are defined by the number of protons in their atomic nucleus, e.g. all xenon atoms have 54 protons. But all elements also have neutrons in their atomic nucleus. Each isotope is characterized by the number of neutrons that it contains, e.g. xenon can have 70, 72, 74, 75, 76, 77, 78, 80, and 82 neutrons, and thus has 9 stable isotopes. The r-process contributes to the abundances of the heaviest four isotopes: 131Xe, 132Xe, 134Xe and 136Xe, and is solely responsible for the heaviest two of those. The s-process contributes to xenon's middle five isotopes: 128Xe, 129Xe, 130Xe, 131Xe, and 132Xe. The lightest two isotopes, 124Xe, and 126Xe, are produced by other processes.

The r-process can typically synthesize the heaviest four isotopes of every heavy element, and the two heaviest isotopes, which are referred to as r-only nuclei, can only be created via the r-process. The r-process abundances peak near atomic weights A = 82 (elements Se, Br and Kr), A = 130 (elements Te, I, and Xe) and A = 196 (elements Os, Ir and Pt).

The r-process entails a succession of rapid neutron captures (hence the name) by one or more heavy seed nuclei, typically beginning with nuclei in the abundance peak centered on 56Fe. The captures must be rapid in the sense that the nuclei must not have time to undergo radioactive decay (typically via β- decay) before another neutron arrives to be captured. This sequence can continue up to the limit of stability of the increasingly neutron-rich nuclei (the neutron drip line) to physically retain neutrons as governed by the short range nuclear force. The r-process therefore must occur in locations where there exist a high density of free neutrons. Early studies theorized that 1024 free neutrons per cm3 would be required, for temperatures about 1GK, in order to match the waiting points, at which no more neutrons can be captured, with the atomic numbers of the abundance peaks for r-process nuclei. This amounts to almost a gram of free neutrons in every cubic centimeter, an astonishing number requiring extreme locations. Traditionally this suggested the material ejected from the reexpanded core of a core-collapse supernova, as part of supernova nucleosynthesis, or decompression of neutron-star matter thrown off by a binary neutron star merger. The relative contributions of these sources to the astrophysical abundance of r-process elements is a matter of ongoing research.A limited r-process-like series of neutron captures occurs to a minor extent in thermonuclear weapon explosions. These led to the discovery of the elements einsteinium (element 99) and fermium (element 100) in nuclear weapon fallout.

The r-process contrasts with the s-process, the other predominant mechanism for the production of heavy elements, which is nucleosynthesis by means of slow captures of neutrons. The s-process primarily occurs within ordinary stars, particularly AGB stars, where the neutron flux is sufficient to cause neutron captures to recur every 10–100 years, much too slow for the r-process, which requires 100 captures per second. The s-process is secondary, meaning that it requires pre-existing heavy isotopes as seed nuclei to be converted into other heavy nuclei by a slow sequence of captures of free neutrons. The r-process scenarios create their own seed nuclei, so they might proceed in massive stars that contain no heavy seed nuclei. Taken together, the r- and s-processes account for almost the entire abundance of chemical elements heavier than iron. The historical challenge has been to locate physical settings appropriate for their time scales.

S-process

The slow neutron-capture process , or s-process is a series of reactions in nuclear astrophysics that occur in stars, particularly AGB stars. The s-process is responsible for the creation (nucleosynthesis) of approximately half the atomic nuclei heavier than iron.

In the s-process, a seed nucleus undergoes neutron capture to form an isotope with one higher atomic mass. If the new isotope is stable, a series of increases in mass can occur, but if it is unstable, then beta decay will occur, producing an element of the next highest atomic number. The process is slow (hence the name) in the sense that there is sufficient time for this radioactive decay to occur before another neutron is captured. A series of these reactions produces stable isotopes by moving along the valley of beta-decay stable isobars in the table of nuclides.

A range of elements and isotopes can be produced by the s-process, because of the intervention of alpha decay steps along the reaction chain. The relative abundances of elements and isotopes produced depends on the source of the neutrons and how their flux changes over time. Each branch of the s-process reaction chain eventually terminates at a cycle involving lead, bismuth, and polonium.

The s-process contrasts with the r-process, in which successive neutron captures are rapid: they happen more quickly than the beta decay can occur. The r-process dominates in environments with higher fluxes of free neutrons; it produces heavier elements and more neutron-rich isotopes than the s-process. Together the two processes account for most of the relative abundance of chemical elements heavier than iron.

Silicon-burning process

In astrophysics, silicon burning is a very brief sequence of nuclear fusion reactions that occur in massive stars with a minimum of about 8-11 solar masses. Silicon burning is the final stage of fusion for massive stars that have run out of the fuels that power them for their long lives in the main sequence on the Hertzsprung-Russell diagram. It follows the previous stages of hydrogen, helium, carbon, neon and oxygen burning processes.

Silicon burning begins when gravitational contraction raises the star's core temperature to 2.7–3.5 billion Kelvin (GK). The exact temperature depends on mass. When a star has completed the silicon-burning phase, no further fusion is possible. The star catastrophically collapses and may explode in what is known as a Type II supernova.

Spectral density estimation

In statistical signal processing, the goal of spectral density estimation (SDE) is to estimate the spectral density (also known as the power spectral density) of a random signal from a sequence of time samples of the signal. Intuitively speaking, the spectral density characterizes the frequency content of the signal. One purpose of estimating the spectral density is to detect any periodicities in the data, by observing peaks at the frequencies corresponding to these periodicities.

Some SDE techniques assume that a signal is composed of a limited (usually small) number of generating frequencies plus noise and seek to find the location and intensity of the generated frequencies. Others make no assumption on the number of components and seek to estimate the whole generating spectrum.

Stellar nucleosynthesis

Stellar nucleosynthesis is the theory explaining the creation (nucleosynthesis) of chemical elements by nuclear fusion reactions between atoms within stars. Stellar nucleosynthesis has occurred continuously since the original creation of hydrogen, helium and lithium during the Big Bang. It is a highly predictive theory that today yields excellent agreement between calculations based upon it and the observed abundances of the elements. It explains why the observed abundances of elements in the universe grow over time and why some elements and their isotopes are much more abundant than others. The theory was initially proposed by Fred Hoyle in 1946, who later refined it in 1954. Further advances were made, especially to nucleosynthesis by neutron capture of the elements heavier than iron, by Margaret Burbidge, Geoffrey Burbidge, William Alfred Fowler and Hoyle in their famous 1957 B2FH paper, which became one of the most heavily cited papers in astrophysics history.

Stars evolve because of changes in their composition (the abundance of their constituent elements) over their lifespans, first by burning hydrogen (main sequence star), then helium (red giant star), and progressively burning higher elements. However, this does not by itself significantly alter the abundances of elements in the universe as the elements are contained within the star. Later in its life, a low-mass star will slowly eject its atmosphere via stellar wind, forming a planetary nebula, while a higher–mass star will eject mass via a sudden catastrophic event called a supernova. The term supernova nucleosynthesis is used to describe the creation of elements during the evolution and explosion of a pre-supernova massive star (12–35 times the mass of the sun). Those massive stars are the most prolific source of new isotopes from carbon (Z = 6) to nickel (Z = 28).

The advanced sequence of burning fuels is driven by gravitational collapse and its associated heating, resulting in the subsequent burning of carbon, oxygen and silicon. However, most of the nucleosynthesis in the mass range A = 28–56 (from silicon to nickel) is actually caused by the upper layers of the star collapsing onto the core, creating a compressional shock wave rebounding outward. The shock front briefly raises temperatures by roughly 50%, thereby causing furious burning for about a second. This final burning in massive stars, called explosive nucleosynthesis or supernova nucleosynthesis, is the final epoch of stellar nucleosynthesis.

A stimulus to the development of the theory of nucleosynthesis was the discovery of variations in the abundances of elements found in the universe. The need for a physical description was already inspired by the relative abundances of isotopes of the chemical elements in the solar system. Those abundances, when plotted on a graph as a function of atomic number of the element, have a jagged sawtooth shape that varies by factors of tens of millions (see history of nucleosynthesis theory). This suggested a natural process that is not random. A second stimulus to understanding the processes of stellar nucleosynthesis occurred during the 20th century, when it was realized that the energy released from nuclear fusion reactions accounted for the longevity of the Sun as a source of heat and light.

Supernova

A supernova ( plural: supernovae or supernovas, abbreviations: SN and SNe) is a transient astronomical event that occurs during the last stellar evolutionary stages of the life of a massive star, whose dramatic and catastrophic destruction is marked by one final, titanic explosion. This causes the sudden appearance of a "new" bright star, before slowly fading from sight over several weeks or months or years.

Supernovae are more energetic than novae. In Latin, nova means "new", referring astronomically to what appears to be a temporary new bright star. Adding the prefix "super-" distinguishes supernovae from ordinary novae, which are far less luminous. The word supernova was coined by Walter Baade and Fritz Zwicky in 1931.

Only three Milky Way, naked-eye supernova events have been observed during the last thousand years, though many have been observed in other galaxies. The most recent directly observed supernova in the Milky Way was Kepler's Supernova in 1604, but the remnants of recent supernovae have also been found. Observations of supernovae in other galaxies suggest they occur on average about three times every century in the Milky Way, and that any galactic supernova would almost certainly be observable with modern astronomical telescopes.

Theoretical studies indicate that most supernovae are triggered by one of two basic mechanisms: the sudden re-ignition of nuclear fusion in a degenerate star or the sudden gravitational collapse of a massive star's core. In the first instance, a degenerate white dwarf may accumulate sufficient material from a binary companion, either through accretion or via a merger, to raise its core temperature enough to trigger runaway nuclear fusion, completely disrupting the star. In the second case, the core of a massive star may undergo sudden gravitational collapse, releasing gravitational potential energy as a supernova. While some observed supernovae are more complex than these two simplified theories, the astrophysical mechanics have been established and accepted by most astronomers for some time.

Supernovae can expel several solar masses of material at speeds up to several percent of the speed of light. This drives an expanding and fast-moving shock wave into the surrounding interstellar medium, sweeping up an expanding shell of gas and dust observed as a supernova remnant. Supernova nucleosynthesis is a major source of elements heavier than nitrogen in the interstellar medium, and the expanding shock waves can directly trigger the formation of new stars. Supernova remnants might be a major source of cosmic rays. Supernovae might produce strong gravitational waves, though, thus far, the gravitational waves detected have been from the merger of black holes and neutron stars, such as those that can be left behind by supernovae.

X-ray burster

X-ray bursters are one class of X-ray binary stars exhibiting periodic and rapid increases in luminosity (typically a factor of 10 or greater) that peak in the X-ray regime of the electromagnetic spectrum. These astrophysical systems are composed of an accreting compact object, and a main sequence companion 'donor' star. A compact object in an X-ray binary system consists of either a neutron star or a black hole; however, with the emission of an X-ray burst, the companion star can immediately be classified as a neutron star, since black holes do not have a surface and all of the accreting material disappears past the event horizon. The donor star's mass falls to the surface of the neutron star where the hydrogen fuses to helium which accumulates until it fuses in a burst, producing X-rays.

The mass of the donor star is used to categorize the system as either a high mass (above 10 solar masses (M☉)) or low mass (less than 1 M☉) X-ray binary, abbreviated as HMXB and LMXB, respectively. X-ray bursters differ observationally from other X-ray transient sources (such as X-ray pulsars and soft X-ray transients), showing a sharp rise time (1 – 10 seconds) followed by spectral softening (a property of cooling black bodies). Individual burst energetics are characterized by an integrated flux of 1032–33 joules, compared to the steady luminosity which is of the order 1032 joules for steady accretion onto a neutron star. As such the ratio α, of the burst flux to the persistent flux, ranges from 10 to 103 but is typically on the order of 100. The X-ray bursts emitted from most of these systems recur on timescales ranging from hours to days, although more extended recurrence times are exhibited in some systems, and weak bursts with recurrence times between 5–20 minutes have yet to be explained but are observed in some less usual cases. The abbreviation XRB can refer either the object (X-ray burster) or the associated emission (X-ray burst). There are two types of XRB's, designated I and II. Type I are far more common than type II, and have a distinctly different cause. Type I are caused by thermonuclear runaway, while type II are caused by gravitational energy release.

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