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In astronomy, metallicity is used to describe the abundance of elements present in an object that are heavier than hydrogen or helium. Most of the physical matter in the Universe is in the form of hydrogen and helium, so astronomers use the word "metals" as a convenient short term for "all elements except hydrogen and helium". This usage is distinct from the usual physical definition of a solid metal. For example, stars and nebulae with relatively high abundances of carbon, nitrogen, oxygen, and neon are called "metal-rich" in astrophysical terms, even though those elements are non-metals in chemistry.

The presence of heavier elements hails from stellar nucleosynthesis, the theory that the majority of elements heavier than hydrogen and helium in the Universe ("metals", hereafter) are formed in the cores of stars as they evolve. Over time, stellar winds and supernovae deposit the metals into the surrounding environment, enriching the interstellar medium and providing recycling materials for the birth of new stars. It follows that older generations of stars, which formed in the metal-poor early 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 astronomer Walter Baade in 1944 to propose the existence of two different populations of stars. These became commonly known as Population I (metal-rich) and Population II (metal-poor) stars. A third stellar population was suggested in 1978, known as Population III stars, which are extremely metal-poor stars and likely the first generation of stars created in the Universe.

Common Methods of Calculation
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 to gas versus metals, or measuring the ratios of the number of atoms of two different elements as compared to the ratios found in the Sun.

Mass Fraction
Stellar composition is usually simply defined by the parameters X, Y and Z. Here X is the fractional percentage of hydrogen, Y is the fractional percentage of helium, and all the remaining chemical elements as the fractional percentage, Z. These values add to unity:


 * $$ X + Y + Z = 1.00 $$

In most stars, nebulae, HII regions, and other astronomical sources, hydrogen and helium are the two dominant elements. The hydrogen mass fraction is generally expressed as $$X\equiv \frac{m_\mathrm{H}}{M}$$ where $$M$$ is the total mass of the system and $$m_\mathrm{H}$$ the total mass of the hydrogen within. Similarly, the helium mass fraction is denoted as $$Y\equiv \frac{m_\mathrm{He}}{M}$$. The remainder of the elements are collectively referred to as 'metals', and the metallicity can be calculated as


 * $$Z = \sum_{i>\mathrm{He}} \frac{m_i}{M} = 1 - X - Y.$$

For the surface of the Sun, these parameters are measured to have the following values: Due to the effects of stellar evolution, neither the initial composition nor the present day bulk composition of the Sun is the same as its present-day surface composition.

Metal to Hydrogen Abundance
The overall stellar metallicity is often defined using the total iron content of the star, as iron is among the easiest to measure with spectral observations in the visible spectrum (even though oxygen is the most abundant heavy element - see Metallicities in HII regions below). The abundance ratio is defined as the logarithm of the ratio of a star's iron content versus it's hydrogen content compared to that of the Sun, and is expressed thus:

$$ [\mathrm{Fe}/\mathrm{H}] = \log_{10}{\left(\frac{N_{\mathrm{Fe}}}{N_{\mathrm{H}}}\right)_\mathrm{star}} - \log_{10}{\left(\frac{N_{\mathrm{Fe}}}{N_{\mathrm{H}}}\right)_\mathrm{sun}} $$

where $$N_{\mathrm{Fe}}$$ and $$N_{\mathrm{H}}$$ are the number of iron and hydrogen atoms per unit of volume, respectively. The unit often used for metallicity is the dex, which is a contraction of 'decimal exponent'. By this formulation, stars with a higher metallicity than the Sun have a positive logarithmic value, whereas those with a lower metallicity than the Sun have a negative value. For example, stars with a [Fe/H] value of +1 have ten times the metallicity of the Sun; conversely, those with a [Fe/H] value of −1 have one-tenth, while those with a [Fe/H] value of 0 have the same metallicity as the Sun, and so on. Young Population I stars have significantly higher iron-to-hydrogen ratios than older Population II stars. Primordial Population III stars are estimated to have metallicities less than −6.0, that is, less than a millionth of the abundance of iron in the Sun.

The same notation is used to express variations in abundances between other individual elements as compared to solar proportions. For example, the notation "[O/Fe]" represents the difference in the logarithm of the star's oxygen abundance versus it's iron content, compared to that of the Sun. In general, a given stellar nucleosynthetic process alters the proportions of only a few elements or isotopes, so a star or gas sample with nonzero [X/Fe] values may be showing the signature of particular nuclear processes.

Photometric Colors
Astronomers can estimate metallicities through measured and calibrated systems that correlate photometric measurements and spectroscopic measurements (see also spectrophotometry). For example, the Johnson UVB filters can be used to detect an ultraviolet (UV) excess in stars, where a larger UV excess indicates a larger presence of metals that absorb the UV radiation, thereby making the star appear 'redder'. The UV excess, δ(U-B), is defined as the difference between a star's U and B band magnitudes, compared to the difference between U and B band magnitudes of metal-rich stars in the Hyades cluster. Unfortunately, δ(U-B) is sensitive to both metallicity and temperature: if two stars are equally metal-rich, but one is cooler than the other, they will likely have different δ(U-B) values (see also line blanketing ). A star's B-V color can be used as an indicator for temperature to mitigate this degeneracy. Furthermore, the UV excess and B-V color can be corrected and normalized to relate the δ(U-B) value to iron abundance.

Other photometric systems that can be used to determine metallicities of certain astrophysical objects include the Strӧmgren system, the Geneva system  , the Washington system  , and the DDO system.

Relationship Between Stellar Metallicity and Planets
A star's metallicity measurement is one parameter that helps determine if a star will have planets and the type of planets, as there is a direct correlation between metallicity and the type of planets a star may have. Measurements have demonstrated the connection between a star's metallicity and gas giant planets, like Jupiter and Saturn. The more metals in a star and thus its planetary system and proplyd, the more likely the system may have gas giant planets and rocky planets. Current models show that the metallicity along with the correct planetary system temperature and distance from the star are key to planet and planetesimal formation. For two stars which have equal age and mass but different metallicity, the less metallic star is bluer. Among stars of the same color, less metallic stars emit more ultraviolet radiation. The Sun, with 8 planets and 5 known dwarf planets, is used as the reference, with a [Fe/H] of 0.00.

HII Regions
Young, massive and hot stars (typically of spectral type O and B) in HII regions emit UV photons that ionize ground state Hydrogen atoms, knocking electrons and protons free; this process is known as photoionization. The free electrons can strike other atoms nearby, exciting bound metallic electrons into a metastable state, which eventually decay back into a ground state, emitting photon energies that correspond to forbidden lines. Through these line transitions, astronomers have developed several observational methods to estimate metal abundances in HII regions where the stronger the forbidden lines in spectroscopic observations, the higher the metallicity. These methods are dependent on one or more of the following: the variety of asymmetrical densities inside HII regions, the varied temperatures of the embedded stars, and/or the electron density within the ionized region.

Theoretically, to determine the total abundance of a single element in an HII region, all transition lines should be observed and summed. However, this can be observationally difficult due to variation in line strength. Some of the most common forbidden lines used to determine metal abundances in HII regions are from oxygen (e.g. [O II] λ = 6300 Å, [O II] λ = (3727, 7318, 7324) Å, and [O III] λ = (4363, 4959, 5007) Å), nitrogen (e.g. [NII] λ = (6548,6584) Å), and sulfur (e.g. [SII] λ = (6717,6731) Å) in the optical spectrum, and the [OIII] λ = (52, 88) μm and [NIII] λ = 57 μm lines in the infrared spectrum. Oxygen is one of the stronger, more abundant lines in HII regions, making it a main target for metallicity estimates within these objects. To measure metal abundances in HII regions using oxygen emission, astronomers use the R23 method, where

$$R_{23} = \frac{[\mathrm{O} II]_{3727 \mathrm{\dot{A}}} + [\mathrm{O} III]_{4959 \mathrm{\dot{A}} + 5007 \mathrm{\dot{A}}}}{\mathrm{H}\beta}$$.

This ratio is well defined through models and observational studies , 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. Similarly, other strong forbidden line ratios can be used; for example, for sulfur where

$$S_{23} = \frac{[\mathrm{S} II]_{6716 \mathrm{\dot{A}} + 6731 \mathrm{\dot{A}}} + [\mathrm{S} III]_{9069 \mathrm{\dot{A}} + 9532 \mathrm{\dot{A}}}}{\mathrm{H}\beta}

$$.

Metal abundances within HII regions are typically less than 1%, with the percentage decreasing on average with distance from the Galactic Center.