Inastronomy,metallicity is theabundance ofelements present in an object that are heavier thanhydrogen andhelium. Most of the normal currently detectable (i.e. non-dark)matter in the universe is either hydrogen or helium, andastronomers use the wordmetals as convenient shorthand forall elements except hydrogen and helium. This word-use is distinct from the conventional chemical or physical definition of ametal as an electrically conducting solid.Stars andnebulae with relatively high abundances of heavier elements are calledmetal-rich when discussing metallicity, even though many of those elements are callednonmetals in chemistry.
Solar spectrum with Fraunhofer lines as it appears visually.
In 1802,William Hyde Wollaston[1] noted the appearance of a number of dark features in the solar spectrum.[2] In 1814,Joseph von Fraunhofer independently rediscovered the lines and began to systematically study and measure theirwavelengths, and they are now calledFraunhofer lines. He mapped over 570 lines, designating the most prominent with the letters A through K and weaker lines with other letters.[3][4][5]
About 45 years later,Gustav Kirchhoff andRobert Bunsen[6] noticed that several Fraunhofer lines coincide with characteristicemission lines identifies in the spectra of heated chemical elements.[7] They inferred that dark lines in the solar spectrum are caused byabsorption bychemical elements in the solar atmosphere.[8] Their observations[9] were in the visible range where the strongest lines come from metals such as sodium, potassium, and iron.[10] In the early work on the chemical composition of the sun the only elements that were detected in spectra were hydrogen and various metals,[11]: 23–24 with the termmetallic frequently used when describing them.[11]: Part 2 In contemporary usage in astronomy all the extra elements beyond just hydrogen and helium are termed metallic.
The presence of heavier elements results from stellar nucleosynthesis, where the majority of elements heavier than hydrogen and helium in the Universe (metals, hereafter) are formed in the cores of stars as theyevolve. Over time,stellar winds andsupernovae deposit the metals into the surrounding environment, enriching theinterstellar medium and providing recycling materials for thebirth of new stars. It follows that older generations of stars, which formed in the metal-poorearly Universe, generally have lower metallicities than those of younger generations, which formed in a more metal-rich Universe.
Observed changes in the chemical abundances of different types of stars, based on the spectral peculiarities that were later attributed to metallicity, led astronomerWalter Baade in 1944 to propose the existence of two differentpopulations of stars.[12]These became commonly known aspopulation I (metal-rich) andpopulation II (metal-poor) stars. A third, earlieststellar population was hypothesized in 1978, known aspopulation III stars.[13][14][15] These "extremely metal-poor" (XMP) stars are theorized to have been the "first-born" stars created in the Universe.
Astronomers use several different methods to describe and approximate metal abundances, depending on the available tools and the object of interest. Some methods include determining the fraction of mass that is attributed togas versus metals, or measuring the ratios of the number of atoms of two different elements as compared to the ratios found in theSun.
Stellar composition is often simply defined by the parametersX,Y, andZ. HereX represents the mass fraction ofhydrogen,Y is the mass fraction ofhelium, andZ is the mass fraction of all the remaining chemical elements. Thus
In moststars,nebulae,H II regions, and other astronomical sources, hydrogen and helium are the two dominant elements. The hydrogen mass fraction is generally expressed as whereM is the total mass of the system, and is the mass of the hydrogen it contains. Similarly, the helium mass fraction is denoted as The remainder of the elements are collectively referred to as "metals", and the mass fraction of metals is calculated as
For the surface of the Sun (symbol), these parameters are measured to have the following values:[16]
Description
Solar value
Hydrogen mass fraction
Helium mass fraction
Metal mass fraction
Due to the effects ofstellar evolution, neither the initial composition nor the present day bulk composition of the Sun is the same as its present-day surface composition.
The overall stellar metallicity is conventionally defined using the total hydrogen content, since its abundance is considered to be relatively constant in the Universe, or theiron content of the star, which has an abundance that is generally linearly increasing in time in the Universe.[17]Hence, iron can be used as a chronological indicator of nucleosynthesis.Iron is relatively easy to measure with spectral observations in the star's spectrum given the large number of iron lines in the star's spectra (even though oxygen is themost abundant heavy element – seemetallicities in H II regions below). The abundance ratio is thecommon logarithm of the ratio of a star's iron abundance compared to that of the Sun and is calculated thus:[18]
where and are the number of iron and hydrogen atoms per unit of volume respectively, is thestandard symbol for the Sun, and for a star (often omitted below). The unit often used for metallicity is thedex, contraction of "decimal exponent".[19] By this formulation, stars with a higher metallicity than the Sun have a positivecommon logarithm, whereas those more dominated by hydrogen have a corresponding negative value. For example, stars with a value of +1 have 10 times the metallicity of the Sun (10+1); conversely, those with a value of −1 have1/10, while those with a value of 0 have the same metallicity as the Sun, and so on.[20]
Young population I stars have significantly higher iron-to-hydrogen ratios than older population II stars. Primordialpopulation III stars are estimated to have metallicity less than −6, a millionth of the abundance of iron in the Sun.[21][22]The same notation is used to express variations in abundances between other individual elements as compared to solar proportions. For example, the notation represents the difference in the logarithm of the star's oxygen abundance versus its iron content compared to that of the Sun. In general, a givenstellar nucleosynthetic process alters the proportions of only a few elements or isotopes, so a star or gas sample with certain values may well be indicative of an associated, studied nuclear process.
Astronomers can estimate metallicities through measured and calibrated systems that correlatephotometric measurements andspectroscopic measurements (see alsoSpectrophotometry). For example, theJohnson UVB filters can be used to detect anultraviolet (UV) excess in stars,[23]where a smaller UV excess indicates a larger presence of metals that absorb the UV radiation, thereby making the star appear "redder".[24][25][26]The UV excess,δ(U−B), is defined as the difference between a star's U and B bandmagnitudes, compared to the difference between U and B band magnitudes of metal-rich stars in theHyades cluster.[27]Unfortunately,δ(U−B) is sensitive to both metallicity andtemperature: If two stars are equally metal-rich, but one is cooler than the other, they will likely have differentδ(U−B) values[27] (see alsoBlanketing effect[28][29]).To help mitigate this degeneracy, a star's B−V color index can be used as an indicator for temperature. Furthermore, the UV excess and B−V index can be corrected to relate theδ(U−B) value to iron abundances.[30][31][32]
Otherphotometric systems that can be used to determine metallicities of certain astrophysical objects include the Strӧmgren system,[33][34]the Geneva system,[35][36] the Washington system,[37][38]and the DDO system.[39][40]
At a given mass and age, a metal-poor star will be slightly warmer.Population II stars' metallicities are roughly1/1000 to1/10 of the Sun's but the group appears cooler thanpopulation I overall, as heavy population II stars have long since died. Above 40 solar masses, metallicity influences how a star will die: Outside thepair-instability window, lower metallicity stars will collapse directly to a black hole, while higher metallicity stars undergo atype Ib/c supernova and may leave aneutron star.
Relationship between stellar metallicity and planets
A star's metallicity measurement is one parameter that helps determine whether a star may have a giantplanet, as there is a direct correlation between metallicity and the presence of a giant planet. Measurements have demonstrated the connection between a star's metallicity andgas giant planets, likeJupiter andSaturn. The more metals in a star and thus itsplanetary system andprotoplanetary disk, the more likely the system may have gas giant planets. Current models show that the metallicity along with the correct planetary system temperature and distance from the star are key to planet andplanetesimal formation. For two stars that have equal age and mass but different metallicity, the less metallic star isbluer. Among stars of the same color, less metallic stars emit more ultraviolet radiation. The Sun, witheight planets and nine consensusdwarf planets, is used as the reference, with a of 0.00.[41][42][43][44][45]
Young, massive and hot stars (typically of spectral typesO andB) inH II regions emitUV photons that ionizeground-state hydrogen atoms, knockingelectrons free; this process is known asphotoionization. The free electrons canstrike other atoms nearby, exciting bound metallic electrons into ametastable state, which eventually decay back into a ground state, emitting photons with energies that correspond toforbidden lines. Through these transitions, astronomers have developed several observational methods to estimate metal abundances in H II regions, where the stronger the forbidden lines in spectroscopic observations, the higher the metallicity.[46][47] These methods are dependent on one or more of the following: the variety of asymmetrical densities inside H II regions, the varied temperatures of the embedded stars, and/or the electron density within the ionized region.[48][49][50][51]
Theoretically, to determine the total abundance of a single element in an H II region, all transition lines should be observed and summed. However, this can be observationally difficult due to variation in line strength.[52][53] Some of the most common forbidden lines used to determine metal abundances in H II regions are fromoxygen (e.g. [OII]λ = (3727, 7318, 7324) Å, and [OIII]λ = (4363, 4959, 5007) Å),nitrogen (e.g. [NII]λ = (5755, 6548, 6584) Å), andsulfur (e.g. [SII]λ = (6717, 6731) Å and [SIII]λ = (6312, 9069, 9531) Å) in theoptical spectrum, and the [OIII]λ = (52, 88) μm and [NIII]λ = 57 μm lines in theinfrared spectrum.Oxygen has some of the stronger, more abundant lines in H II regions, making it a main target for metallicity estimates within these objects. To calculate metal abundances in H II regions using oxygenflux measurements, astronomers often use theR23 method, in which
where is the sum of the fluxes from oxygenemission lines measured at therest frameλ = (3727, 4959 and 5007) Å wavelengths, divided by the flux from theBalmer series Hβ emission line at the rest frameλ = 4861 Å wavelength.[54]This ratio is well defined through models and observational studies,[55][56][57] but caution should be taken, as the ratio is often degenerate, providing both a low and high metallicity solution, which can be broken with additional line measurements.[58]Similarly, other strong forbidden line ratios can be used, e.g. for sulfur, where[59]
Metal abundances within H II regions are typically less than 1%, with the percentage decreasing on average with distance from theGalactic Center.[52][60][61][62][63]
Gustav Kirchhoff (1859)"Ueber die Fraunhofer'schen Linien" (On Fraunhofer's lines),Monatsbericht der Königlichen Preussische Akademie der Wissenschaften zu Berlin (Monthly report of the Royal Prussian Academy of Sciences in Berlin), 662–665.
Gustav Kirchhoff (1859)"Ueber das Sonnenspektrum" (On the sun's spectrum),Verhandlungen des naturhistorisch-medizinischen Vereins zu Heidelberg (Proceedings of the Natural History / Medical Association in Heidelberg),1 (7) : 251–255.
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