In astronomy, metallicity is the abundance of elements present in an object that are heavier than hydrogen and helium. Most of the normal currently detectable (i.e. non-dark) matter in the universe is either hydrogen or helium, and astronomers use the word metals as convenient shorthand for all elements except hydrogen and helium. This word-use is distinct from the conventional chemical or physical definition of a metal as an electrically conducting element. Stars and nebulae with relatively high abundances of heavier elements are called metal-rich when discussing metallicity, even though many of those elements are called nonmetals in chemistry.

In 1802, William Hyde Wollaston noted the appearance of a number of dark features in the solar spectrum. In 1814, Joseph von Fraunhofer independently rediscovered the lines and began to systematically study and measure their wavelengths, and they are now called Fraunhofer lines. He mapped over 570 lines, designating the most prominent with the letters A through K and weaker lines with other letters.

About 45 years later, Gustav Kirchhoff and Robert Bunsen noticed that several Fraunhofer lines coincide with characteristic emission lines identified in the spectra of heated chemical elements. They inferred that dark lines in the solar spectrum are caused by absorption by chemical elements in the solar atmosphere. Their observations were in the visible range where the strongest lines come from metals such as sodium, potassium, and iron. In the early work on the chemical composition of the sun the only elements that were detected in spectra were hydrogen and various metals, with the term metallic frequently used when describing them. 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 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, earliest stellar population was hypothesized in 1978, known as population III stars. 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 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.

Stellar composition is often simply defined by the parameters X, Y, and Z. Here X represents the mass fraction of hydrogen, Y is the mass fraction of helium, and Z is the mass fraction of all the remaining chemical elements. Thus

$${\displaystyle X+Y+Z=1}$$