A variable star is a star whose brightness as seen from Earth fluctuates. This variation may be caused by a change in emitted light or by something blocking the light, so variable stars are classified as either: Intrinsic variables, whose luminosity changes. Extrinsic variables, whose apparent changes in brightness are due to changes in the amount of their light that can reach Earth. Many most, stars have at least some variation in luminosity: the energy output of our Sun, for example, varies by about 0.1% over an 11-year solar cycle. An ancient Egyptian calendar of lucky and unlucky days composed some 3,200 years ago may be the oldest preserved historical document of the discovery of a variable star, the eclipsing binary Algol. Of the modern astronomers, the first variable star was identified in 1638 when Johannes Holwarda noticed that Omicron Ceti pulsated in a cycle taking 11 months; this discovery, combined with supernovae observed in 1572 and 1604, proved that the starry sky was not eternally invariable as Aristotle and other ancient philosophers had taught.
In this way, the discovery of variable stars contributed to the astronomical revolution of the sixteenth and early seventeenth centuries. The second variable star to be described was the eclipsing variable Algol, by Geminiano Montanari in 1669. Chi Cygni was identified in 1686 by G. Kirch R Hydrae in 1704 by G. D. Maraldi. By 1786 ten variable stars were known. John Goodricke himself discovered Beta Lyrae. Since 1850 the number of known variable stars has increased especially after 1890 when it became possible to identify variable stars by means of photography; the latest edition of the General Catalogue of Variable Stars lists more than 46,000 variable stars in the Milky Way, as well as 10,000 in other galaxies, over 10,000'suspected' variables. The most common kinds of variability involve changes in brightness, but other types of variability occur, in particular changes in the spectrum. By combining light curve data with observed spectral changes, astronomers are able to explain why a particular star is variable.
Variable stars are analysed using photometry, spectrophotometry and spectroscopy. Measurements of their changes in brightness can be plotted to produce light curves. For regular variables, the period of variation and its amplitude can be well established. Peak brightnesses in the light curve are known as maxima. Amateur astronomers can do useful scientific study of variable stars by visually comparing the star with other stars within the same telescopic field of view of which the magnitudes are known and constant. By estimating the variable's magnitude and noting the time of observation a visual lightcurve can be constructed; the American Association of Variable Star Observers collects such observations from participants around the world and shares the data with the scientific community. From the light curve the following data are derived: are the brightness variations periodical, irregular, or unique? What is the period of the brightness fluctuations? What is the shape of the light curve? From the spectrum the following data are derived: what kind of star is it: what is its temperature, its luminosity class? is it a single star, or a binary? does the spectrum change with time?
Changes in brightness may depend on the part of the spectrum, observed if the wavelengths of spectral lines are shifted this points to movements strong magnetic fields on the star betray themselves in the spectrum abnormal emission or absorption lines may be indication of a hot stellar atmosphere, or gas clouds surrounding the star. In few cases it is possible to make pictures of a stellar disk; these may show darker spots on its surface. Combining light curves with spectral data gives a clue as to the changes that occur in a variable star. For example, evidence for a pulsating star is found in its shifting spectrum because its surface periodically moves toward and away from us, with the same frequency as its changing brightness. About two-thirds of all variable stars appear to be pulsating. In the 1930s astronomer Arthur Stanley Eddington showed that the mathematical equations that describe the interior of a star may lead to instabilities that cause a star to pulsate; the most common type of instability is related to oscillations in the degree of ionization in outer, convective layers of the star.
Suppose the star is in the swelling phase. Its outer layers expand; because of the decreasing temperature the degree of ionization decreases. This makes the gas more transparent, thus makes it easier for the star to radiate its energy; this in turn will make the star start to contract. As the gas is thereby compressed, it is heated and the degree of ionization again increases. Thi
The parsec is a unit of length used to measure large distances to astronomical objects outside the Solar System. A parsec is defined as the distance at which one astronomical unit subtends an angle of one arcsecond, which corresponds to 648000/π astronomical units. One parsec is equal to 31 trillion kilometres or 19 trillion miles; the nearest star, Proxima Centauri, is about 1.3 parsecs from the Sun. Most of the stars visible to the unaided eye in the night sky are within 500 parsecs of the Sun; the parsec unit was first suggested in 1913 by the British astronomer Herbert Hall Turner. Named as a portmanteau of the parallax of one arcsecond, it was defined to make calculations of astronomical distances from only their raw observational data quick and easy for astronomers. For this reason, it is the unit preferred in astronomy and astrophysics, though the light-year remains prominent in popular science texts and common usage. Although parsecs are used for the shorter distances within the Milky Way, multiples of parsecs are required for the larger scales in the universe, including kiloparsecs for the more distant objects within and around the Milky Way, megaparsecs for mid-distance galaxies, gigaparsecs for many quasars and the most distant galaxies.
In August 2015, the IAU passed Resolution B2, which, as part of the definition of a standardized absolute and apparent bolometric magnitude scale, mentioned an existing explicit definition of the parsec as 648000/π astronomical units, or 3.08567758149137×1016 metres. This corresponds to the small-angle definition of the parsec found in many contemporary astronomical references; the parsec is defined as being equal to the length of the longer leg of an elongated imaginary right triangle in space. The two dimensions on which this triangle is based are its shorter leg, of length one astronomical unit, the subtended angle of the vertex opposite that leg, measuring one arc second. Applying the rules of trigonometry to these two values, the unit length of the other leg of the triangle can be derived. One of the oldest methods used by astronomers to calculate the distance to a star is to record the difference in angle between two measurements of the position of the star in the sky; the first measurement is taken from the Earth on one side of the Sun, the second is taken half a year when the Earth is on the opposite side of the Sun.
The distance between the two positions of the Earth when the two measurements were taken is twice the distance between the Earth and the Sun. The difference in angle between the two measurements is twice the parallax angle, formed by lines from the Sun and Earth to the star at the distant vertex; the distance to the star could be calculated using trigonometry. The first successful published direct measurements of an object at interstellar distances were undertaken by German astronomer Friedrich Wilhelm Bessel in 1838, who used this approach to calculate the 3.5-parsec distance of 61 Cygni. The parallax of a star is defined as half of the angular distance that a star appears to move relative to the celestial sphere as Earth orbits the Sun. Equivalently, it is the subtended angle, from that star's perspective, of the semimajor axis of the Earth's orbit; the star, the Sun and the Earth form the corners of an imaginary right triangle in space: the right angle is the corner at the Sun, the corner at the star is the parallax angle.
The length of the opposite side to the parallax angle is the distance from the Earth to the Sun (defined as one astronomical unit, the length of the adjacent side gives the distance from the sun to the star. Therefore, given a measurement of the parallax angle, along with the rules of trigonometry, the distance from the Sun to the star can be found. A parsec is defined as the length of the side adjacent to the vertex occupied by a star whose parallax angle is one arcsecond; the use of the parsec as a unit of distance follows from Bessel's method, because the distance in parsecs can be computed as the reciprocal of the parallax angle in arcseconds. No trigonometric functions are required in this relationship because the small angles involved mean that the approximate solution of the skinny triangle can be applied. Though it may have been used before, the term parsec was first mentioned in an astronomical publication in 1913. Astronomer Royal Frank Watson Dyson expressed his concern for the need of a name for that unit of distance.
He proposed the name astron, but mentioned that Carl Charlier had suggested siriometer and Herbert Hall Turner had proposed parsec. It was Turner's proposal. In the diagram above, S represents the Sun, E the Earth at one point in its orbit, thus the distance ES is one astronomical unit. The angle SDE is one arcsecond so by definition D is a point in space at a distance of one parsec from the Sun. Through trigonometry, the distance SD is calculated as follows: S D = E S tan 1 ″ S D ≈ E S 1 ″ = 1 au 1 60 × 60 × π
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 and nebulae with high abundances of carbon, nitrogen and neon are called "metal-rich" in astrophysical terms 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 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 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 attributed to metallicity, led astronomer Walter Baade in 1944 to propose the existence of two different populations of stars. These became known as Population I and Population II stars. A third stellar population was introduced in 1978, known as Population III stars; these metal-poor stars were theorised 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, 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 simply defined by the parameters X, Y and Z. Here X is the mass fraction of hydrogen, Y is the mass fraction of helium, Z is the mass fraction of all the remaining chemical elements.
Thus X + Y + Z = 1.00. In most stars, nebulae, H II regions, other astronomical sources and helium are the two dominant elements; the hydrogen mass fraction is expressed as X ≡ m H / M, where M is the total mass of the system, m H is the fractional mass of the hydrogen it contains. The helium mass fraction is denoted as Y ≡ m He / M; the remainder of the elements are collectively referred to as "metals", the metallicity—the mass fraction of elements heavier than helium—can be calculated as Z = ∑ i > He 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; the overall stellar metallicity is defined using the total iron content of the star, as iron is among the easiest to measure with spectral observations in the visible spectrum. The abundance ratio is defined as the logarithm of the ratio of a star's iron abundance compared to that of the Sun and is expressed thus: = log 10 star − log 10 sun, where N Fe and N H are the number of iron and hydrogen atoms per unit of volume respectively.
The unit used for metallicity is the dex, 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 value of +1 have 10 times the metallicity of the Sun. Young Population I stars have higher iron-to-hydrogen ratios than older Population II stars. Primordial Population III stars are estimated to have a metallicity of 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 the individual elements as compared to solar proportions. For example, the notati
A giant star is a star with larger radius and luminosity than a main-sequence star of the same surface temperature. They lie above the main sequence on the Hertzsprung–Russell diagram and correspond to luminosity classes II and III; the terms giant and dwarf were coined for stars of quite different luminosity despite similar temperature or spectral type by Ejnar Hertzsprung about 1905. Giant stars have radii up to a few hundred times the Sun and luminosities between 10 and a few thousand times that of the Sun. Stars still more luminous than giants are referred to as hypergiants. A hot, luminous main-sequence star may be referred to as a giant, but any main-sequence star is properly called a dwarf no matter how large and luminous it is. A star becomes a giant after all the hydrogen available for fusion at its core has been depleted and, as a result, leaves the main sequence; the behaviour of a post-main-sequence star depends on its mass. For a star with a mass above about 0.25 solar masses, once the core is depleted of hydrogen it contracts and heats up so that hydrogen starts to fuse in a shell around the core.
The portion of the star outside the shell expands and cools, but with only a small increase in luminosity, the star becomes a subgiant. The inert helium core continues to grow and increase temperature as it accretes helium from the shell, but in stars up to about 10-12 M☉ it does not become hot enough to start helium burning. Instead, after just a few million years the core reaches the Schönberg–Chandrasekhar limit collapses, may become degenerate; this causes the outer layers to expand further and generates a strong convective zone that brings heavy elements to the surface in a process called the first dredge-up. This strong convection increases the transport of energy to the surface, the luminosity increases and the star moves onto the red-giant branch where it will stably burn hydrogen in a shell for a substantial fraction of its entire life; the core continues to gain mass and increase in temperature, whereas there is some mass loss in the outer layers. § 5.9. If the star's mass, when on the main sequence, was below 0.4 M☉, it will never reach the central temperatures necessary to fuse helium.
P. 169. It will therefore remain a hydrogen-fusing red giant until it runs out of hydrogen, at which point it will become a helium white dwarf. § 4.1, 6.1. According to stellar evolution theory, no star of such low mass can have evolved to that stage within the age of the Universe. In stars above about 0.4 M☉ the core temperature reaches 108 K and helium will begin to fuse to carbon and oxygen in the core by the triple-alpha process.§ 5.9, chapter 6. When the core is degenerate helium fusion begins explosively, but most of the energy goes into lifting the degeneracy and the core becomes convective; the energy generated by helium fusion reduces the pressure in the surrounding hydrogen-burning shell, which reduces its energy-generation rate. The overall luminosity of the star decreases, its outer envelope contracts again, the star moves from the red-giant branch to the horizontal branch. Chapter 6; when the core helium is exhausted, a star with up to about 8 M☉ has a carbon–oxygen core that becomes degenerate and starts helium burning in a shell.
As with the earlier collapse of the helium core, this starts convection in the outer layers, triggers a second dredge-up, causes a dramatic increase in size and luminosity. This is the asymptotic giant branch analogous to the red-giant branch but more luminous, with a hydrogen-burning shell contributing most of the energy. Stars only remain on the AGB for around a million years, becoming unstable until they exhaust their fuel, go through a planetary nebula phase, become a carbon–oxygen white dwarf. § 7.1–7.4. Main-sequence stars with masses above about 12 M☉ are very luminous and they move horizontally across the HR diagram when they leave the main sequence becoming blue giants before they expand further into blue supergiants, they start core-helium burning before the core becomes degenerate and develop smoothly into red supergiants without a strong increase in luminosity. At this stage they have comparable luminosities to bright AGB stars although they have much higher masses, but will further increase in luminosity as they burn heavier elements and become a supernova.
Stars in the 8-12 M☉ range have somewhat intermediate properties and have been called super-AGB stars. They follow the tracks of lighter stars through RGB, HB, AGB phases, but are massive enough to initiate core carbon burning and some neon burning, they form oxygen–magnesium–neon cores, which may collapse in an electron-capture supernova, or they may leave behind an oxygen–neon white dwarf. O class main sequence stars are highly luminous; the giant phase for such stars is a brief phase of increased size and luminosity before developing a supergiant spectral luminosity class. Type O giants may be more than a hundred thousand times as luminous as the sun, brighter than many supergiants. Classification is complex and difficult with small differences between luminosity classes and a continuous range of intermediate forms; the most massive stars develop giant or supergiant spectral features while still burning hydrogen in their cores, due to mixing of heavy elements to the surface and high luminosity which produces a powerful stellar wind and causes the star's atmosphere to expand.
A star whose initial mass is less than 0.25 M☉ will not become a giant star at all. For most of th
Astrometry is the branch of astronomy that involves precise measurements of the positions and movements of stars and other celestial bodies. The information obtained by astrometric measurements provides information on the kinematics and physical origin of the Solar System and our galaxy, the Milky Way; the history of astrometry is linked to the history of star catalogues, which gave astronomers reference points for objects in the sky so they could track their movements. This can be dated back to Hipparchus, who around 190 BC used the catalogue of his predecessors Timocharis and Aristillus to discover Earth's precession. In doing so, he developed the brightness scale still in use today. Hipparchus compiled a catalogue with their positions. Hipparchus's successor, included a catalogue of 1,022 stars in his work the Almagest, giving their location and brightness. In the 10th century, Abd al-Rahman al-Sufi carried out observations on the stars and described their positions and star color. Ibn Yunus observed more than 10,000 entries for the Sun's position for many years using a large astrolabe with a diameter of nearly 1.4 metres.
His observations on eclipses were still used centuries in Simon Newcomb's investigations on the motion of the Moon, while his other observations of the motions of the planets Jupiter and Saturn inspired Laplace's Obliquity of the Ecliptic and Inequalities of Jupiter and Saturn. In the 15th century, the Timurid astronomer Ulugh Beg compiled the Zij-i-Sultani, in which he catalogued 1,019 stars. Like the earlier catalogs of Hipparchus and Ptolemy, Ulugh Beg's catalogue is estimated to have been precise to within 20 minutes of arc. In the 16th century, Tycho Brahe used improved instruments, including large mural instruments, to measure star positions more than with a precision of 15–35 arcsec. Taqi al-Din measured the right ascension of the stars at the Constantinople Observatory of Taqi ad-Din using the "observational clock" he invented; when telescopes became commonplace, setting circles sped measurements James Bradley first tried to measure stellar parallaxes in 1729. The stellar movement proved too insignificant for his telescope, but he instead discovered the aberration of light and the nutation of the Earth's axis.
His cataloguing of 3222 stars was refined in 1807 by Friedrich Bessel, the father of modern astrometry. He made the first measurement of stellar parallax: 0.3 arcsec for the binary star 61 Cygni. Being difficult to measure, only about 60 stellar parallaxes had been obtained by the end of the 19th century by use of the filar micrometer. Astrographs using astronomical photographic plates sped the process in the early 20th century. Automated plate-measuring machines and more sophisticated computer technology of the 1960s allowed more efficient compilation of star catalogues. In the 1980s, charge-coupled devices replaced photographic plates and reduced optical uncertainties to one milliarcsecond; this technology made astrometry less expensive. In 1989, the European Space Agency's Hipparcos satellite took astrometry into orbit, where it could be less affected by mechanical forces of the Earth and optical distortions from its atmosphere. Operated from 1989 to 1993, Hipparcos measured large and small angles on the sky with much greater precision than any previous optical telescopes.
During its 4-year run, the positions and proper motions of 118,218 stars were determined with an unprecedented degree of accuracy. A new "Tycho catalog" drew together a database of 1,058,332 to within 20-30 mas. Additional catalogues were compiled for the 23,882 double/multiple stars and 11,597 variable stars analyzed during the Hipparcos mission. Today, the catalogue most used is USNO-B1.0, an all-sky catalogue that tracks proper motions, positions and other characteristics for over one billion stellar objects. During the past 50 years, 7,435 Schmidt camera plates were used to complete several sky surveys that make the data in USNO-B1.0 accurate to within 0.2 arcsec. Apart from the fundamental function of providing astronomers with a reference frame to report their observations in, astrometry is fundamental for fields like celestial mechanics, stellar dynamics and galactic astronomy. In observational astronomy, astrometric techniques help identify stellar objects by their unique motions, it is instrumental for keeping time, in that UTC is the atomic time synchronized to Earth's rotation by means of exact astronomical observations.
Astrometry is an important step in the cosmic distance ladder because it establishes parallax distance estimates for stars in the Milky Way. Astrometry has been used to support claims of extrasolar planet detection by measuring the displacement the proposed planets cause in their parent star's apparent position on the sky, due to their mutual orbit around the center of mass of the system. Astrometry is more accurate in space missions that are not affected by the distorting effects of the Earth's atmosphere. NASA's planned Space Interferometry Mission was to utilize astrometric techniques to detect terrestrial planets orbiting 200 or so of the nearest solar-type stars; the European Space Agency's Gaia Mission, launched in 2013, applies astrometric techniques in its stellar census. In addition to the detection of exoplanets, it can be used to determine their mass. Astrometric measurements are used by astrophysicists to constrain certain models in celestial mechanics. By measuring the velocities of pulsars, it is possible to put a limit on the asymmetry of supernova explosions.
Right ascension is the angular distance of a particular point measured eastward along the celestial equator from the Sun at the March equinox to the point above the earth in question. When paired with declination, these astronomical coordinates specify the direction of a point on the celestial sphere in the equatorial coordinate system. An old term, right ascension refers to the ascension, or the point on the celestial equator that rises with any celestial object as seen from Earth's equator, where the celestial equator intersects the horizon at a right angle, it contrasts with oblique ascension, the point on the celestial equator that rises with any celestial object as seen from most latitudes on Earth, where the celestial equator intersects the horizon at an oblique angle. Right ascension is the celestial equivalent of terrestrial longitude. Both right ascension and longitude measure an angle from a primary direction on an equator. Right ascension is measured from the Sun at the March equinox i.e. the First Point of Aries, the place on the celestial sphere where the Sun crosses the celestial equator from south to north at the March equinox and is located in the constellation Pisces.
Right ascension is measured continuously in a full circle from that alignment of Earth and Sun in space, that equinox, the measurement increasing towards the east. As seen from Earth, objects noted to have 12h RA are longest visible at the March equinox. On those dates at midnight, such objects will reach their highest point. How high depends on their declination. Any units of angular measure could have been chosen for right ascension, but it is customarily measured in hours and seconds, with 24h being equivalent to a full circle. Astronomers have chosen this unit to measure right ascension because they measure a star's location by timing its passage through the highest point in the sky as the Earth rotates; the line which passes through the highest point in the sky, called the meridian, is the projection of a longitude line onto the celestial sphere. Since a complete circle contains 24h of right ascension or 360°, 1/24 of a circle is measured as 1h of right ascension, or 15°. A full circle, measured in right-ascension units, contains 24 × 60 × 60 = 86400s, or 24 × 60 = 1440m, or 24h.
Because right ascensions are measured in hours, they can be used to time the positions of objects in the sky. For example, if a star with RA = 1h 30m 00s is at its meridian a star with RA = 20h 00m 00s will be on the/at its meridian 18.5 sidereal hours later. Sidereal hour angle, used in celestial navigation, is similar to right ascension, but increases westward rather than eastward. Measured in degrees, it is the complement of right ascension with respect to 24h, it is important not to confuse sidereal hour angle with the astronomical concept of hour angle, which measures angular distance of an object westward from the local meridian. The Earth's axis rotates westward about the poles of the ecliptic, completing one cycle in about 26,000 years; this movement, known as precession, causes the coordinates of stationary celestial objects to change continuously, if rather slowly. Therefore, equatorial coordinates are inherently relative to the year of their observation, astronomers specify them with reference to a particular year, known as an epoch.
Coordinates from different epochs must be mathematically rotated to match each other, or to match a standard epoch. Right ascension for "fixed stars" near the ecliptic and equator increases by about 3.05 seconds per year on average, or 5.1 minutes per century, but for fixed stars further from the ecliptic the rate of change can be anything from negative infinity to positive infinity. The right ascension of Polaris is increasing quickly; the North Ecliptic Pole in Draco and the South Ecliptic Pole in Dorado are always at right ascension 18h and 6h respectively. The used standard epoch is J2000.0, January 1, 2000 at 12:00 TT. The prefix "J" indicates. Prior to J2000.0, astronomers used the successive Besselian epochs B1875.0, B1900.0, B1950.0. The concept of right ascension has been known at least as far back as Hipparchus who measured stars in equatorial coordinates in the 2nd century BC, but Hipparchus and his successors made their star catalogs in ecliptic coordinates, the use of RA was limited to special cases.
With the invention of the telescope, it became possible for astronomers to observe celestial objects in greater detail, provided that the telescope could be kept pointed at the object for a period of time. The easiest way to do, to use an equatorial mount, which allows the telescope to be aligned with one of its two pivots parallel to the Earth's axis. A motorized clock drive is used with an equatorial mount to cancel out the Earth's rotation; as the equatorial mount became adopted for observation, the equatorial coordinate system, which includes right ascension, was adopted at the same time for simplicity. Equatorial mounts could be pointed at objects with known right ascension and declination by the use of setting circles; the first star catalog to use right ascen
Astronomical spectroscopy is the study of astronomy using the techniques of spectroscopy to measure the spectrum of electromagnetic radiation, including visible light and radio, which radiates from stars and other celestial objects. A stellar spectrum can reveal many properties of stars, such as their chemical composition, density, distance and relative motion using Doppler shift measurements. Spectroscopy is used to study the physical properties of many other types of celestial objects such as planets, nebulae and active galactic nuclei. Astronomical spectroscopy is used to measure three major bands of radiation: visible spectrum, X-ray. While all spectroscopy looks at specific areas of the spectrum, different methods are required to acquire the signal depending on the frequency. Ozone and molecular oxygen absorb light with wavelengths under 300 nm, meaning that X-ray and ultraviolet spectroscopy require the use of a satellite telescope or rocket mounted detectors. Radio signals have much longer wavelengths than optical signals, require the use of antennas or radio dishes.
Infrared light is absorbed by atmospheric water and carbon dioxide, so while the equipment is similar to that used in optical spectroscopy, satellites are required to record much of the infrared spectrum. Physicists have been looking at the solar spectrum since Isaac Newton first used a simple prism to observe the refractive properties of light. In the early 1800s Joseph von Fraunhofer used his skills as a glass maker to create pure prisms, which allowed him to observe 574 dark lines in a continuous spectrum. Soon after this, he combined telescope and prism to observe the spectrum of Venus, the Moon and various stars such as Betelgeuse; the resolution of a prism is limited by its size. This issue was resolved in the early 1900s with the development of high-quality reflection gratings by J. S. Plaskett at the Dominion Observatory in Ottawa, Canada. Light striking a mirror will reflect at the same angle, however a small portion of the light will be refracted at a different angle. By creating a "blazed" grating which utilizes a large number of parallel mirrors, the small portion of light can be focused and visualized.
These new spectroscopes were more detailed than a prism, required less light, could be focused on a specific region of the spectrum by tilting the grating. The limitation to a blazed grating is the width of the mirrors, which can only be ground a finite amount before focus is lost. In order to overcome this limitation holographic gratings were developed. Volume phase holographic gratings use a thin film of dichromated gelatin on a glass surface, subsequently exposed to a wave pattern created by an interferometer; this wave pattern sets up a reflection pattern similar to the blazed gratings but utilizing Bragg diffraction, a process where the angle of reflection is dependent on the arrangement of the atoms in the gelatin. The holographic gratings can have up to 6000 lines/mm and can be up to twice as efficient in collecting light as blazed gratings; because they are sealed between two sheets of glass, the holographic gratings are versatile lasting decades before needing replacement. Light dispersed by the grating or prism in a spectrograph can be recorded by a detector.
Photographic plates were used to record spectra until electronic detectors were developed, today optical spectrographs most employ charge-coupled devices. The wavelength scale of a spectrum can be calibrated by observing the spectrum of emission lines of known wavelength from a gas-discharge lamp; the flux scale of a spectrum can be calibrated as a function of wavelength by comparison with an observation of a standard star with corrections for atmospheric absorption of light. Radio astronomy was founded with the work of Karl Jansky in the early 1930s, while working for Bell Labs, he built a radio antenna to look at potential sources of interference for transatlantic radio transmissions. One of the sources of noise discovered came not from Earth, but from the center of the Milky Way, in the constellation Sagittarius. In 1942, JS Hey captured the sun's radio frequency using military radar receivers. Radio spectroscopy started with the discovery of the 21-centimeter H I line in 1951. Radio interferometry was pioneered in 1946, when Joseph Lade Pawsey, Ruby Payne-Scott and Lindsay McCready used a single antenna atop a sea cliff to observe 200 MHz solar radiation.
Two incident beams, one directly from the sun and the other reflected from the sea surface, generated the necessary interference. The first multi-receiver interferometer was built in the same year by Martin Vonberg. In 1960, Ryle and Antony Hewish published the technique of aperture synthesis to analyze interferometer data; the aperture synthesis process, which involves autocorrelating and discrete Fourier transforming the incoming signal, recovers both the spatial and frequency variation in flux. The result is a 3D image. For this work and Hewish were jointly awarded the 1974 Nobel Prize in Physics. Newton used a prism to split white light into a spectrum of color, Fraunhofer's high-quality prisms allowed scientists to see dark lines of an unknown origin. In the 1850s, Gustav Kirchhoff and Robert Bunsen described the phenomena behind these dark lines