* Your assessment is very important for improving the workof artificial intelligence, which forms the content of this project
Download Observational Constraints The Nebular Hypothesis
Observational astronomy wikipedia , lookup
Kepler (spacecraft) wikipedia , lookup
Astronomical unit wikipedia , lookup
History of astronomy wikipedia , lookup
Space Interferometry Mission wikipedia , lookup
Spitzer Space Telescope wikipedia , lookup
Circumstellar habitable zone wikipedia , lookup
Astrobiology wikipedia , lookup
Aquarius (constellation) wikipedia , lookup
Rare Earth hypothesis wikipedia , lookup
Planets beyond Neptune wikipedia , lookup
Beta Pictoris wikipedia , lookup
Planets in astrology wikipedia , lookup
Dwarf planet wikipedia , lookup
Extraterrestrial life wikipedia , lookup
Directed panspermia wikipedia , lookup
Satellite system (astronomy) wikipedia , lookup
Star formation wikipedia , lookup
IAU definition of planet wikipedia , lookup
Accretion disk wikipedia , lookup
Definition of planet wikipedia , lookup
Solar System wikipedia , lookup
Late Heavy Bombardment wikipedia , lookup
Exoplanetology wikipedia , lookup
Planetary habitability wikipedia , lookup
Timeline of astronomy wikipedia , lookup
Formation and evolution of the Solar System wikipedia , lookup
History of Solar System formation and evolution hypotheses wikipedia , lookup
The Sun and Planets Lecture Notes 2. Spring Semester 2017 Prof Dr Ravit Helled Lecture 2 Formation of the Solar System Observational Constraints Observational Constraints The following are the observed features of the Solar System that any formation theory has to reproduce. We call these features “observational constraints” because they constrain our formation theories. Age — The age of the Solar System is 4.567 billion years (Gyr). Orbit shape — The orbits of the planets are ellipses with the Sun at one focus. The orbits of the planets are all more or less in the same plane. Orbit direction — The planets orbits are prograde (i.e., in the same direction as the Sun’s rotation). Orbit distance — The planets’ orbits lie within 30 AU of the Sun and the spacing between the planets is large. Mass distribution — The planets only account for < 0.2% of the Solar System’s total mass. The Nebular Hypothesis Roughly 4.6 billion years ago, a great cloud of gas and dust (i.e., a nebula) began to collapse when gravitational forces became stronger than the cloud’s internal gas pressure. The cloud was composed mostly of hydrogen and helium, with smaller fractions of heavier elements. As the cloud collapsed, three things happened: 1. The cloud begins to rotate faster (conservation of angular momentum), 2. The cloud flattens into a disk, 3. The interior of the cloud heats up. At the center, the temperature continues to rise until it reaches ∼1 million degrees Kelvin. This is the temperature required to trigger thermonuclear fusion (i.e., the star “turns on”). 1 Protoplanetary Disk The disk of gas and dust that surrounds a young star is called a “protoplanetary disk”. This is the disk of material that planets form in. The temperature at any given location in the disk determines whether certain materials at that location will be solid of gaseous. The disk is hotter closer to the Sun, so the composition of the disk will change depending on the distance from the center. In the inner disk, the gas and dust are strongly heated by the Sun. In this inner region, ices and gasses cannot condense to form solids. Particles that condense in the inner region are mainly silicates and iron compounds (i.e., rocks and metals). In the cold outer disk, ices and gasses are able to condense to form solids. The boundary that separates the two regions is know as the “snowline”. When we discuss the snowline, we usually mean the H2 O snowline, however each material has a different condensation temperature and therefore a different snowline (e.g., iron compounds have a much higher condensation temperature and therefore their snowline will be much closer to the Sun). Planet Composition Metal is found in the cores of planets, mostly in the form of iron (Fe) and nickel (Ni). Rock: the most common rocks are silicates. Silicates are oxides of silicon (Si), magnesium (Mg), and Aluminium (Al). Ices are also referred to as “volatiles”. The most common ices are water (H2 O), carbon dioxide (CO2 ), methane (CH4 ), and ammonia (NH3 ). Gas in planets is predominately hydrogen (H) and helium (He). The origin of these gases is the nebula (i.e., cloud of gas and dust) that the Solar System formed from. 2 Planet Formation Terrestrial planets: 1. Small dust grains grow into larger—but still relatively small—asteroid-like bodies called planetesimals. 2. Planetesimals repeated crash into each other, resulting in increasingly large planetesimals. Some of these objects grow large enough to be called protoplanets. 3. As the protoplanets grow to their final size, they slowly clear out the Solar System of most smaller bodies (this process is called disk dispersal). Giant planets: There are two suggested models to explain how giant planets are formed: core accretion and disk instability. Core accretion: Planet agglomerates from dust and attracts gas envelope. In this scenario, giant planets should have non-stellar compositions and massive cores. Core accretion proceeds through the following three steps: 1. Accretion of dust particles and planetesimals results in a core of a few Earth masses with a low-mass gaseous envelope. 2. Further accretion of gas and solids: the envelope grows faster than the core until the cross-over mass is reached. 3. Runaway gas accretion with relatively small accretion of solids. Disk instability: Clump of gas inside the protoplanetary disk collapses to form a giant planet. In this scenario giant planets should have stellar compositions and no cores. The time scale for giant planet formation via disk instability is roughly 1000 years. The Toomre parameter, Q, is a way to measure when and where the disk will become unstable and potentially lead to giant planet formation: Q= cs Ω πGσg where cs is the sound speed, Ω is the angular velocity, G is the gravitational constant, and σg is the surface density. For the ideal case of an infinitesimally thin disk, Q > 1 is stable and Q < 1 is unstable. The disk instability process is most efficient as large radial distances. 3 Giant Planets In the Solar System: Jupiter, Saturn, Uranus, and Neptune. In the Solar System we can identify a few trends in physical and orbital characteristics. 1. Giant planets exist at large radial distances (> 5 AU). 2. The mass of the giant planets decreases with orbital distance. 3. Heavy element enrichment increases with decreasing mass. Outside of the Solar System: 1. While there are many observed extrasolar giant planets, they still appear to be less common than small (terrestrial) planets. An up-to-date catalogue of observed extrasolar planets can be found at http://exoplanets.eu/. 2. Giant planets exist at very small orbital distances (e.g., the record for the shortest orbital period is held by PSR J1719-1438 b, a Jupiter mass planet orbiting its host star at 0.004 AU). 3. The giant planet occurrence rate is ∼5–20%. The occurrence rate is a measure of how many stars host giant planets (e.g., if we look at 100 stars and find giant planets orbiting 20 of them, the occurrence rate is 20%). The variation between 5% and 20% reflects the fact that the occurrence rate varies with stellar mass and metallicity. 4. Observations of protoplanetary disks around other stars suggest that the typical disk lifetime is < 10 million years and the typical disk mass is 0.01–0.1 M . This tells us that giant planets have to form in less than 10 Myr and from a disk of material that is 1% to 10% the mass of the Sun). 5. Transiting giant planets contain 0–100 M⊕ of heavy elements. 6. Giant planets exist at very large radial distances. We know this thanks to a technique called direct imaging, whereby astronomers are able take pictures of planets around other stars. 7. Giant planets have been observed around M dwarf stars (< 0.6 M ) and metal-poor stars. 4 Stellar Composition & Metallicity Stellar composition, as determined by spectroscopy, is usually defined by the parameters X, Y, and Z. In this system, X is the fractional percentage of hydrogen, Y is the fractional percentage of helium, and Z is the fractional percentage of all remaining elements (i.e., the “metals”). X, Y, and Z must add up to exactly one: X + Y + Z = 1.0 The overall stellar metallicity is defined using the total iron (Fe) content of the star relative to its hydrogen (H) content. Stellar metallicity is denoted as [Fe/H]. Note that iron is not the most abundant heavy element—oxygen is—but it is the easiest for astronomers to measure using spectroscopy. The metallicity is commonly defined as the logarithm of the ratio of a star’s iron abundance compared to that of the Sun: [Fe/H] = log10 NF e NH − log10 star NF e NH Sun where NF e and NH are the number of iron and hydrogen atoms per unit of volume, respectively. These values we can determine from spectroscopy. 5