Gaia is a space observatory of the European Space Agency, launched in 2013 and expected to operate until c. 2022. The spacecraft is designed for astrometry: measuring the positions and motions of stars with unprecedented precision; the mission aims to construct by far the largest and most precise 3D space catalog made, totalling 1 billion astronomical objects stars, but planets, comets and quasars among others. The spacecraft has monitored each of its target objects about 70 times over the first five years of the mission to study the precise position and motion of each target, will keep doing so; the spacecraft has enough micro-propulsion fuel to operate until about November 2024. As its detectors are not degrading as fast as expected, the mission could therefore be extended; the Gaia targets represent 1% of the Milky Way population with all stars brighter than magnitude 20 in a broad photometric band that covers most of the visual range. Additionally, Gaia is expected to detect thousands to tens of thousands of Jupiter-sized exoplanets beyond the Solar System, 500,000 quasars outside our galaxy and tens of thousands of new asteroids and comets within the Solar System.
Gaia will create a precise three-dimensional map of astronomical objects throughout the Milky Way and map their motions, which encode the origin and subsequent evolution of the Milky Way. The spectrophotometric measurements will provide the detailed physical properties of all stars observed, characterizing their luminosity, effective temperature and elemental composition; this massive stellar census will provide the basic observational data to analyze a wide range of important questions related to the origin and evolutionary history of our galaxy. The successor to the Hipparcos mission, Gaia is part of ESA's Horizon 2000+ long-term scientific program. Gaia was launched on 19 December 2013 by Arianespace using a Soyuz ST-B/Fregat-MT rocket flying from Kourou in French Guiana; the spacecraft operates in a Lissajous orbit around the Sun–Earth L2 Lagrangian point. The Gaia space telescope has its roots in ESA's Hipparcos mission, its mission was proposed in October 1993 by Lennart Lindegren and Michael Perryman in response to a call for proposals for ESA's Horizon Plus long-term scientific programme.
It was adopted by ESA's Science Programme Committee as cornerstone mission number 6 on 13 October 2000, the B2 phase of the project was authorised on 9 February 2006, with EADS Astrium taking responsibility for the hardware. The name "Gaia" was derived as an acronym for Global Astrometric Interferometer for Astrophysics; this reflected the optical technique of interferometry, planned for use on the spacecraft. While the working method evolved during studies and the acronym is no longer applicable, the name Gaia remained to provide continuity with the project; the total cost of the mission is around €740 million, including the manufacture and ground operations. Gaia was completed two years behind schedule and 16% above its initial budget due to the difficulties encountered in polishing Gaia's ten mirrors and assembling and testing the focal plane camera system; the Gaia space mission has the following objectives: To determine the intrinsic luminosity of a star requires knowledge of its distance.
One of the few ways to achieve this without physical assumptions is through the star's parallax. Ground-based observations would not measure such parallaxes with sufficient precision due to the effects of the atmosphere and instrumental biases. For instance, Cepheid variables are used as standard candles to measure distances to galaxies, but the accuracy in their own distance measurement is poor. Thus, quantities depending on them, such as the speed of expansion of the universe, remain inaccurate. Measuring their distances has a great impact on the understanding of the other galaxies and thus the whole cosmos. Observations of the faintest objects will provide a more complete view of the stellar luminosity function. Gaia will observe 1 billion stars and other bodies, representing 1% of such bodies in the Milky Way galaxy. All objects up to a certain magnitude must be measured in order to have unbiased samples. To permit a better understanding of the more rapid stages of stellar evolution; this has to be achieved by detailed examination and re-examination of a great number of objects over a long period of operation.
Observing a large number of objects in the galaxy is important to understand the dynamics of our galaxy. Measuring the astrometric and kinematic properties of a star is necessary in order to understand the various stellar populations the most distant. In order to achieve these objectives, Gaia has these goals: Determine the position and annual proper motion of 1 billion stars with an accuracy of about 20 microarcseconds at 15 mag, 200 µas at 20 mag. Determine the positions of stars at a magnitude of V = 10 down to a precision of 7 μas—this is equivalent to measuring the position to within the diameter of a hair from 1000 km away—between 12 and 25 μas down to V = 15, between 100 and 300 μas to V = 20, depending on the colour of the star; the distance to about 20 million stars will thus be measured with a precision of 1% or better, about 200 million distances will be measured to better than 10%. Distances accurate to 10% will be achieved as far away as the Galactic Centre, 30,000 light-years away.
Measure the tangential speed of 40 million stars to a precision of better
The Sun is the star at the center of the Solar System. It is a nearly perfect sphere of hot plasma, with internal convective motion that generates a magnetic field via a dynamo process, it is by far the most important source of energy for life on Earth. Its diameter is about 1.39 million kilometers, or 109 times that of Earth, its mass is about 330,000 times that of Earth. It accounts for about 99.86% of the total mass of the Solar System. Three quarters of the Sun's mass consists of hydrogen; the Sun is a G-type main-sequence star based on its spectral class. As such, it is informally and not accurately referred to as a yellow dwarf, it formed 4.6 billion years ago from the gravitational collapse of matter within a region of a large molecular cloud. Most of this matter gathered in the center, whereas the rest flattened into an orbiting disk that became the Solar System; the central mass became so hot and dense that it initiated nuclear fusion in its core. It is thought that all stars form by this process.
The Sun is middle-aged. It fuses about 600 million tons of hydrogen into helium every second, converting 4 million tons of matter into energy every second as a result; this energy, which can take between 10,000 and 170,000 years to escape from its core, is the source of the Sun's light and heat. In about 5 billion years, when hydrogen fusion in its core has diminished to the point at which the Sun is no longer in hydrostatic equilibrium, its core will undergo a marked increase in density and temperature while its outer layers expand to become a red giant, it is calculated that the Sun will become sufficiently large to engulf the current orbits of Mercury and Venus, render Earth uninhabitable. After this, it will shed its outer layers and become a dense type of cooling star known as a white dwarf, no longer produce energy by fusion, but still glow and give off heat from its previous fusion; the enormous effect of the Sun on Earth has been recognized since prehistoric times, the Sun has been regarded by some cultures as a deity.
The synodic rotation of Earth and its orbit around the Sun are the basis of solar calendars, one of, the predominant calendar in use today. The English proper name Sun may be related to south. Cognates to English sun appear in other Germanic languages, including Old Frisian sunne, Old Saxon sunna, Middle Dutch sonne, modern Dutch zon, Old High German sunna, modern German Sonne, Old Norse sunna, Gothic sunnō. All Germanic terms for the Sun stem from Proto-Germanic *sunnōn; the Latin name for the Sun, Sol, is not used in everyday English. Sol is used by planetary astronomers to refer to the duration of a solar day on another planet, such as Mars; the related word solar is the usual adjectival term used for the Sun, in terms such as solar day, solar eclipse, Solar System. A mean Earth solar day is 24 hours, whereas a mean Martian'sol' is 24 hours, 39 minutes, 35.244 seconds. The English weekday name Sunday stems from Old English and is a result of a Germanic interpretation of Latin dies solis, itself a translation of the Greek ἡμέρα ἡλίου.
The Sun is a G-type main-sequence star. The Sun has an absolute magnitude of +4.83, estimated to be brighter than about 85% of the stars in the Milky Way, most of which are red dwarfs. The Sun is heavy-element-rich, star; the formation of the Sun may have been triggered by shockwaves from more nearby supernovae. This is suggested by a high abundance of heavy elements in the Solar System, such as gold and uranium, relative to the abundances of these elements in so-called Population II, heavy-element-poor, stars; the heavy elements could most plausibly have been produced by endothermic nuclear reactions during a supernova, or by transmutation through neutron absorption within a massive second-generation star. The Sun is by far the brightest object in the Earth's sky, with an apparent magnitude of −26.74. This is about 13 billion times brighter than the next brightest star, which has an apparent magnitude of −1.46. The mean distance of the Sun's center to Earth's center is 1 astronomical unit, though the distance varies as Earth moves from perihelion in January to aphelion in July.
At this average distance, light travels from the Sun's horizon to Earth's horizon in about 8 minutes and 19 seconds, while light from the closest points of the Sun and Earth takes about two seconds less. The energy of this sunlight supports all life on Earth by photosynthesis, drives Earth's climate and weather; the Sun does not have a definite boundary, but its density decreases exponentially with increasing height above the photosphere. For the purpose of measurement, the Sun's radius is considered to be the distance from its center to the edge of the photosphere, the apparent visible surface of the Sun. By this measure, the Sun is a near-perfect sphere with an oblateness estimated at about 9 millionths, which means that its polar diameter differs from its equatorial diameter by only 10 kilometres; the tidal effect of the planets is weak and does not affect the shape of the Sun. The Sun rotates faster at its equator than at its poles; this differential rotation is caused by convective motion
R Aquarii is a variable star in the constellation Aquarius. R Aquarii is a symbiotic star believed to contain a white dwarf and a Mira-type variable in a binary system; the orbital period is 44 years. The main Mira-type star is a red giant, varies in brightness by a factor of several hundred and with a period of more than a year, it has a distance of about 200 parsec, is one of the nearest symbiotic stars and a well known jet source. The two components have been resolved at a separation of 55 mas. By its gravitational pull, the white dwarf draws in material from the red giant and ejects some of the surplus in weird loops to form the nebula seen in the linked image; the whole system appears reddened because it is situated in a dusty region of space, its blue light is absorbed before reaching Earth. The nebula around R Aquarii is known as Cederblad 211, it is possible that the nebula is the remnant of a nova-like outburst, which may have been observed by Japanese astronomers, in the year 930 AD. It is reasonably small and dominated by its central star.
Visual observations are rare. The central region of the jet shows an ejection that took place around 190 years ago, as well as much younger structures; the giant primary star is a Mira variable, a star that pulsates and changes temperature, leading to large visual brightness changes. The total range of 5.2 - 12.4 is a variation of 750 times in brightness, from a naked eye star to one beyond the range of binoculars. The pulsations occur every 390 days but are not regular. Weird loop image in R Aquarii R Aquarii on WikiSky: DSS2, SDSS, GALEX, IRAS, Hydrogen α, X-Ray, Sky Map and images NASA Astronomy Picture of the Day: Symbiotic Star System R Aquarii AAVSO Variable Star of the Month. R Aquarii: Summer 2003 R Aquarii Visual observation of a symbiotic star
In astronomy, luminosity is the total amount of energy emitted per unit of time by a star, galaxy, or other astronomical object. As a term for energy emitted per unit time, luminosity is synonymous with power. In SI units luminosity is measured in joules per second or watts. Values for luminosity are given in the terms of the luminosity of the Sun, L⊙. Luminosity can be given in terms of the astronomical magnitude system: the absolute bolometric magnitude of an object is a logarithmic measure of its total energy emission rate, while absolute magnitude is a logarithmic measure of the luminosity within some specific wavelength range or filter band. In contrast, the term brightness in astronomy is used to refer to an object's apparent brightness: that is, how bright an object appears to an observer. Apparent brightness depends on both the luminosity of the object and the distance between the object and observer, on any absorption of light along the path from object to observer. Apparent magnitude is a logarithmic measure of apparent brightness.
The distance determined by luminosity measures can be somewhat ambiguous, is thus sometimes called the luminosity distance. In astronomy, luminosity is the amount of electromagnetic energy; when not qualified, the term "luminosity" means bolometric luminosity, measured either in the SI units, watts, or in terms of solar luminosities. A bolometer is the instrument used to measure radiant energy over a wide band by absorption and measurement of heating. A star radiates neutrinos, which carry off some energy, contributing to the star's total luminosity; the IAU has defined a nominal solar luminosity of 3.828×1026 W to promote publication of consistent and comparable values in units of the solar luminosity. While bolometers do exist, they cannot be used to measure the apparent brightness of a star because they are insufficiently sensitive across the electromagnetic spectrum and because most wavelengths do not reach the surface of the Earth. In practice bolometric magnitudes are measured by taking measurements at certain wavelengths and constructing a model of the total spectrum, most to match those measurements.
In some cases, the process of estimation is extreme, with luminosities being calculated when less than 1% of the energy output is observed, for example with a hot Wolf-Rayet star observed only in the infra-red. Bolometric luminosities can be calculated using a bolometric correction to a luminosity in a particular passband; the term luminosity is used in relation to particular passbands such as a visual luminosity of K-band luminosity. These are not luminosities in the strict sense of an absolute measure of radiated power, but absolute magnitudes defined for a given filter in a photometric system. Several different photometric systems exist; some such as the UBV or Johnson system are defined against photometric standard stars, while others such as the AB system are defined in terms of a spectral flux density. A star's luminosity can be determined from two stellar characteristics: size and effective temperature; the former is represented in terms of solar radii, R⊙, while the latter is represented in kelvins, but in most cases neither can be measured directly.
To determine a star's radius, two other metrics are needed: the star's angular diameter and its distance from Earth. Both can be measured with great accuracy in certain cases, with cool supergiants having large angular diameters, some cool evolved stars having masers in their atmospheres that can be used to measure the parallax using VLBI. However, for most stars the angular diameter or parallax, or both, are far below our ability to measure with any certainty. Since the effective temperature is a number that represents the temperature of a black body that would reproduce the luminosity, it cannot be measured directly, but it can be estimated from the spectrum. An alternative way to measure stellar luminosity is to measure the star's apparent brightness and distance. A third component needed to derive the luminosity is the degree of interstellar extinction, present, a condition that arises because of gas and dust present in the interstellar medium, the Earth's atmosphere, circumstellar matter.
One of astronomy's central challenges in determining a star's luminosity is to derive accurate measurements for each of these components, without which an accurate luminosity figure remains elusive. Extinction can only be measured directly if the actual and observed luminosities are both known, but it can be estimated from the observed colour of a star, using models of the expected level of reddening from the interstellar medium. In the current system of stellar classification, stars are grouped according to temperature, with the massive young and energetic Class O stars boasting temperatures in excess of 30,000 K while the less massive older Class M stars exhibit temperatures less than 3,500 K; because luminosity is proportional to temperature to the fourth power, the large variation in stellar temperatures produces an vaster variation in stellar luminosity. Because the luminosity depends on a high power of the stellar mass, high mass luminous stars have much shorter lifetimes; the most luminous stars are always young stars, no more than a few million years for the most extreme.
In the Hertzsprung–Russell diagram, the x-axis represents temperature or spectral type while the y-axis represents luminosity or magnitude. The vast majority of stars are found along the main sequence with blue Class O stars found at the top left of the chart while red Class M stars fall to the bottom right. Certain stars like Deneb and Betelgeuse are
A red giant is a luminous giant star of low or intermediate mass in a late phase of stellar evolution. The outer atmosphere is inflated and tenuous, making the radius large and the surface temperature around 5,000 K or lower; the appearance of the red giant is from yellow-orange to red, including the spectral types K and M, but class S stars and most carbon stars. The most common red giants are stars on the red-giant branch that are still fusing hydrogen into helium in a shell surrounding an inert helium core. Other red giants are the red-clump stars in the cool half of the horizontal branch, fusing helium into carbon in their cores via the triple-alpha process. Red giants are stars that have exhausted the supply of hydrogen in their cores and have begun thermonuclear fusion of hydrogen in a shell surrounding the core, they have radii tens to hundreds of times larger than that of the Sun. However, their outer envelope is lower in temperature, giving them a reddish-orange hue. Despite the lower energy density of their envelope, red giants are many times more luminous than the Sun because of their great size.
Red-giant-branch stars have luminosities up to nearly three thousand times that of the Sun, spectral types of K or M, have surface temperatures of 3,000–4,000 K, radii up to about 200 times the Sun. Stars on the horizontal branch are hotter, with only a small range of luminosities around 75 L☉. Asymptotic-giant-branch stars range from similar luminosities as the brighter stars of the red giant branch, up to several times more luminous at the end of the thermal pulsing phase. Among the asymptotic-giant-branch stars belong the carbon stars of type C-N and late C-R, produced when carbon and other elements are convected to the surface in what is called a dredge-up; the first dredge-up occurs during hydrogen shell burning on the red-giant branch, but does not produce a large carbon abundance at the surface. The second, sometimes third, dredge up occurs during helium shell burning on the asymptotic-giant branch and convects carbon to the surface in sufficiently massive stars; the stellar limb of a red giant is not defined, contrary to their depiction in many illustrations.
Rather, due to the low mass density of the envelope, such stars lack a well-defined photosphere, the body of the star transitions into a'corona'. The coolest red giants have complex spectra, with molecular lines, emission features, sometimes masers from thermally pulsing AGB stars. Another noteworthy feature of red giants is that, unlike Sun-like stars whose photospheres have a large number of small convection cells, red-giant photospheres, as well as those of red supergiants, have just a few large cells, the features of which cause the variations of brightness so common on both types of stars. Red giants are evolved from main-sequence stars with masses in the range from about 0.3 M☉ to around 8 M☉. When a star forms from a collapsing molecular cloud in the interstellar medium, it contains hydrogen and helium, with trace amounts of "metals"; these elements are all uniformly mixed throughout the star. The star reaches the main sequence when the core reaches a temperature high enough to begin fusing hydrogen and establishes hydrostatic equilibrium.
Over its main sequence life, the star converts the hydrogen in the core into helium. For the Sun, the main-sequence lifetime is 10 billion years. More-massive stars burn disproportionately faster and so have a shorter lifetime than less massive stars; when the star exhausts the hydrogen fuel in its core, nuclear reactions can no longer continue and so the core begins to contract due to its own gravity. This brings additional hydrogen into a zone where the temperature and pressure are adequate to cause fusion to resume in a shell around the core; the outer layers of the star expand thus beginning the red-giant phase of the star's life. As the star expands, the energy produced in the burning shell of the star is spread over a much larger surface area, resulting in a lower surface temperature and a shift in the star's visible light output towards the red—hence it becomes a red giant. At this time, the star is said to be ascending the red-giant branch of the Hertzsprung–Russell diagram; the evolutionary path the star takes as it moves along the red-giant branch, which ends with the complete collapse of the core, depends on the mass of the star.
For the Sun and stars of less than about 2 M☉ the core will become dense enough that electron degeneracy pressure will prevent it from collapsing further. Once the core is degenerate, it will continue to heat until it reaches a temperature of 108 K, hot enough to begin fusing helium to carbon via the triple-alpha process. Once the degenerate core reaches this temperature, the entire core will begin helium fusion nearly in a so-called helium flash. In more-massive stars, the collapsing core will reach 108 K before it is dense enough to be degenerate, so helium fusion will begin much more smoothly, produce no helium flash; the core helium fusing phase of a star's life is called the horizontal branch in metal-poor stars, so named because these stars lie on a nearly horizontal line in the H–R diagram of many star clusters. Metal-rich helium-fusing stars instead lie on the so-called red clump in the H–R diagram. An analogous process oc
Helium is a chemical element with symbol He and atomic number 2. It is a colorless, tasteless, non-toxic, monatomic gas, the first in the noble gas group in the periodic table, its boiling point is the lowest among all the elements. After hydrogen, helium is the second lightest and second most abundant element in the observable universe, being present at about 24% of the total elemental mass, more than 12 times the mass of all the heavier elements combined, its abundance is similar in Jupiter. This is due to the high nuclear binding energy of helium-4 with respect to the next three elements after helium; this helium-4 binding energy accounts for why it is a product of both nuclear fusion and radioactive decay. Most helium in the universe is helium-4, the vast majority of, formed during the Big Bang. Large amounts of new helium are being created by nuclear fusion of hydrogen in stars. Helium is named for the Greek Titan of the Sun, Helios, it was first detected as an unknown yellow spectral line signature in sunlight during a solar eclipse in 1868 by Georges Rayet, Captain C. T. Haig, Norman R. Pogson, Lieutenant John Herschel, was subsequently confirmed by French astronomer Jules Janssen.
Janssen is jointly credited with detecting the element along with Norman Lockyer. Janssen recorded the helium spectral line during the solar eclipse of 1868 while Lockyer observed it from Britain. Lockyer was the first to propose; the formal discovery of the element was made in 1895 by two Swedish chemists, Per Teodor Cleve and Nils Abraham Langlet, who found helium emanating from the uranium ore cleveite. In 1903, large reserves of helium were found in natural gas fields in parts of the United States, by far the largest supplier of the gas today. Liquid helium is used in cryogenics in the cooling of superconducting magnets, with the main commercial application being in MRI scanners. Helium's other industrial uses—as a pressurizing and purge gas, as a protective atmosphere for arc welding and in processes such as growing crystals to make silicon wafers—account for half of the gas produced. A well-known but minor use is as a lifting gas in airships; as with any gas whose density differs from that of air, inhaling a small volume of helium temporarily changes the timbre and quality of the human voice.
In scientific research, the behavior of the two fluid phases of helium-4 is important to researchers studying quantum mechanics and to those looking at the phenomena, such as superconductivity, produced in matter near absolute zero. On Earth it is rare—5.2 ppm by volume in the atmosphere. Most terrestrial helium present today is created by the natural radioactive decay of heavy radioactive elements, as the alpha particles emitted by such decays consist of helium-4 nuclei; this radiogenic helium is trapped with natural gas in concentrations as great as 7% by volume, from which it is extracted commercially by a low-temperature separation process called fractional distillation. Terrestrial helium—a non-renewable resource, because once released into the atmosphere it escapes into space—was thought to be in short supply. However, recent studies suggest that helium produced deep in the earth by radioactive decay can collect in natural gas reserves in larger than expected quantities, in some cases having been released by volcanic activity.
The first evidence of helium was observed on August 18, 1868, as a bright yellow line with a wavelength of 587.49 nanometers in the spectrum of the chromosphere of the Sun. The line was detected by French astronomer Jules Janssen during a total solar eclipse in Guntur, India; this line was assumed to be sodium. On October 20 of the same year, English astronomer Norman Lockyer observed a yellow line in the solar spectrum, which he named the D3 because it was near the known D1 and D2 Fraunhofer line lines of sodium, he concluded. Lockyer and English chemist Edward Frankland named the element with the Greek word for the Sun, ἥλιος. In 1881, Italian physicist Luigi Palmieri detected helium on Earth for the first time through its D3 spectral line, when he analyzed a material, sublimated during a recent eruption of Mount Vesuvius. On March 26, 1895, Scottish chemist Sir William Ramsay isolated helium on Earth by treating the mineral cleveite with mineral acids. Ramsay was looking for argon but, after separating nitrogen and oxygen from the gas liberated by sulfuric acid, he noticed a bright yellow line that matched the D3 line observed in the spectrum of the Sun.
These samples were identified as helium by Lockyer and British physicist William Crookes. It was independently isolated from cleveite in the same year by chemists Per Teodor Cleve and Abraham Langlet in Uppsala, who collected enough of the gas to determine its atomic weight. Helium was isolated by the American geochemist William Francis Hillebrand prior to Ramsay's discovery when he noticed unusual spectral lines while testing a sample of the mineral uraninite. Hillebrand, attributed the lines to nitrogen, his letter of congratulations to Ramsay offers an interesting case of discovery and near-discovery in science. In 1907, Ernest Rutherford and Thomas Royds demonstrated that alpha particles are helium nuclei by allowing the particles to penetrate the thin glass wall of
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