The Interstellar Medium Gas and Dust between the stars…. Sometimes Interstellar Matter is bright and colorful... Emission Nebula – glowing gas Sometimes Interstellar Matter is “Dark”….. Dark Clouds The dark regions are due to dust obscuring the background starlight. The Milky Way: View towards the Galactic Center The dark regions are due to dust obscuring the backggground starlight. Dust Properties: Dust grains are tiny! ~100-1000 nm (compara ble to v is ible li gh t) Dust absorption + scattering = extinction. Dust absorbs/scatters blue light more than red light. ⇒ Backdlihkground light gets “reddene d” as it passes through interstellar dust. Light can only be absorbed or scattered by particles having diameters comppgarable to or larger than the wavelen gggth of that light! Dust absorbs/scatters blue light more than red light. Why is the sky BLUE? Atmospppheric particles are about the size of the blue wavelength of light. The Sun’s blue light is scattered in all directions – we are seeing this scattered light in the sky. The Milky Way: View towards the Galactic Center Near Infrared (2 microns) Visible Light Interstellar reddening occurs because longer wavelengths pass through dusty clouds more easily. We can “see” through dust clouds in red, infrared and radio light. Stars embedded in a Dark Cloud… This is one of the best ways to study star formation. visible light (few stars) infrared light (many stars!) How do we learn about DUST in the ISM (InterStellar Medium)? Consider a dust cloud between a star and the Earth: • Dust reduces the intensity of light, with larger extinctions at shorter wavelengths. • Stellar absorption lines are st ill p rese nt and can be used to determine spectral type of star . • If the spectral type of a star is known, its luminosity, temperature and true spectral shape are known. • Comparing the true stellar spectrum to the observed spectrum tells us how much dust the light has traveled thru. Dust Temperatures: Thermal (Blackbody) Emission. Warm dust grains emit like blackbodies, with spectra peaking in the infrared. The peak wavelength reveals the temperature (Wien’s Law). Horsehead Nebula (Orion) visible light In visible light , the dust causes obscuration (dark clouds) . But in the far infrared……. Horsehead Nebula (Orion) Far infrared light The dust is emitting at these wavelengths. How do we learn about GAS in the ISM? Orion Hot gas clouds (near stars) emit line radiation…. Different colors are caused by different emission lines. Red: Hα 6563A (3→2) Green: O+2 5007A Eagle Nebula (M16, in Orion) How do we learn about gas in the ISM? Emission lines tell us about the composition, temperature, density and motion of the interstellar gas. Emission Line Nebulae: Also called HII regions (ionized H = H+) “Photoionized” by hot stars. Element abundances siiltSimilar to Sun. Temperatures: ~8000K Densities: ~10-100 ions/cm3 Lagoon Nebula (M8) Average ISM Characteristics: ((ymostly outside dense clouds) Gas Temperatures: a few K to a few hundred K (water freezes at 273 K) Gas Density: 1 atom / cm3 (ranges from 0.01 up to 106 in dense clouds) (best laboratory Vacuums are 10,000 atoms / cm3) Dust Density: 10-12 particles / cm3 ((pthis is 1 particle in a box 328 feet on a side ) Summary: How do we measure the ISM? Dust extinction: dims and reddens light from distant stars. DtDust em iiission: warm dus t em its lik e a bl ackb od y (i nf rared) . Emission-Line Nebulae: Atoms in hot gas emit spectral lines. (near hot stars only!) Absorption line spectra : Atoms in cool gas absorb in lines . (in front of stars only!) Problem: Most interstellar gas is cool, and we can’t count on having bright background stars, so… How can we measure it? Measuring cool gas: 21-cm Radiation Bu mping atomic H from gro und (n =1) to e xcited (n =2 ,3 ,…) state requires substantial energy (thousands of K). But the n=1 level is split, depending on spin of proton and electron. A spin flip changes the energy very slightly, emitting a long-wavelength photon (21 cm). The spins are either parallel or anti-parallel (quantum mechanics). slightly excited state true ground state Measuring cool gas: 21-cm Radiation The Galaxy is nearly transparent to 21 cm line radiation. Atomic Hydrogen 21 cm Visible Light Near Infrared (2 microns) The visible light is from stars, but obscured. The IR light is also from stars , much less obscured . The 21 cm line radiation is from cool interstellar gas (H atoms). Measuring cool gas: Interstellar Molecules Dense dark clouds (>104 cm-3 and not near a star) contain mostly molecules, not atoms. H2 is by far the most abundant. Other molecules: CO (carbon monoxide) HCN (hydrogen cyanide) NH3 (ammonia) H2O(O (water ) CH3OH (methyl alcohol) H2CO (formaldehyde) Some molecules are much more complex… MlMolecu les appear only in dense dar k c loud s… Wh y ? (p. 286) Molecules emit/absorb light via rotation, vibration or electronic transitions. A transition from rapid to slower rotation in formaldehyde emits one radio photon. VibtilttiltitiVibrational or rotational transitions are most usef flbul, because they occur at IR – radio wavelengths, …allowing us to “see” into dense dark molecular clouds . This formaldehyde “map” shows where the cool dense gas lies…. Trifid Nebula (M20) Contours indicate the amount of formaldehyde. Different contours correspond to different trans iti ons. Density ~105 cm-3 Temperature ~50K 4 pc Why do we care about the ISM? 1) It affects our measurements of distant stars and galaxies. 2) It is an important component of galaxies. 3) Stars and Galaxies form out of gas & dust. Eagle Nebula (M16, in Orion) Eagle Nebula (M16, in Orion) Dense dkldark clou ds are t he birthplaces of stars… Why do stars form? → Gravity. Every particle in the universe is attracted to every other particle. They are most attracted to their nearest neighbors because of 1/d2. Interstellar clouds are pulled inward by their own “weight.” Wha t coun terac ts this ten dency for “co llapse?” → Heat = Motion. (Remember: Temperature is a measure of motion.) The gravitational force of a few atoms on each other is small compared to the effect of heat (random motion of the atoms) so they do not bind together into a lasting clump of matter. How many atoms (how much total mass) is required for the collapse of an ISM cloud? That depends on Temperature! For T = 100 K, ~1057 atoms are needed (~1 solar mass)! Stage 1: Cloud collapse and Fragmentation (2 million years) Clouds are non-uniform. The densest pockets collapse first, leading to fragmentation. Initial cloud size: 1 – 100 pc Total initial mass: ~10 – 1000 solar mass Stars form in groups - notit isol ltdated Stage 2: Continuing Collapse of each Fragment (3000 years) Gravitational collapse should increase the kinetic energy of the c loud → motion = heat. But,,g the added heat goes into “excitin g” atoms and molecules, which quickly radiate that energy away. We call this “radiative cooling.” Stage 2: Continuing Collapse of each Fragment. Fragment masses are a bit larger than the stars they will create . Bok Globules Stage 3: Fragmentation Ceases, a Protostar is Born (10000 years) Central regions of the fragment become opaque to their own radiation. Trapped radiation means poorer cooling. The central temperature becomes much higher than the surroundings. The dense opaque inner region is called a Protostar. Protostars have a “surface” (like stars) called the photosphere. Stage 4: Protostar I (1 million years) The Pro tos tar ( opaque core) continues to contract, …and gain mass as material “rains” onto it. Contraction leads to continued heating…. Building pressure (from heat) starts to slow the collapse. High luminosity caused by release o f grav itati onal energy Stage 5: Protostar II (10 million years) Contraction proceeds slowly as the internal temperature and pressure increase… These gaseous disks make planets (sometimes?). Stage 6: A Star is Born! (30 million years) The central temperature reaches 107 K. Nuclear reactions begin… A Star is Born! The radius and temppgygeratures are still slightly greater than the Sun. Contraction continues very slowly as the stars settles into equilibrium… (gravity versus internal pressure). Stage 7: The Main Sequence (10 billion years) The newly formed star takes ~ 30 m illion more years to reach the Main Sequence… (for a solar mass star). The star then spends ~10 billion years on the MS, burning H in its core… Stars of other masses form similarly… …populating different parts ofthf the ma in sequence. But the evolution times are shorter for more massive stars. Star Clusters – the result of fragmentation •Contain stars of different masses,,p that land on different parts of the main sequence. •Excellent “laboratories” for testing theories of stellar evolution. Pleiades (open cluster) Failed Stars: Brown Dwarfs Some fragments do not have enough mass to become stars. They reach equilibrium (pressure vs . gravity) but their central temperatures remain too low for nuclear burning. Other examples: gaseous planets like Jupiter… The minimum mass for a star is ~0.08 solar masses. Brown Dwarfs never reach the main sequence, they just fade away… Gliese 623 + brown dwarf.