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Reading Assignment Hubble, 1947, PASP, 59, 153 (part of HW #1 due today) Questions: What are the major advances that Hubble predicts the 200-inch telescope will offer to astronomers? Was he correct? This week at Astro 3303 • HW #1 is due today; please pass it in. • After it is graded, you can include it in your portfolio • HW #2 is now posted; it is due next Tuesday. It includes use of TOPCAT to investigate some datasets. You may work on the homework together but: 1. Be sure to name your collaborators 2. Hand in your own copy of the answers • Pick up PE #3 Future event: Remote Observing with Arecibo Feb 22 @10pm What is the purpose of a telescope? 1. A telescope acts like a light bucket, to gather photons. • The bigger a telescope is, the more photons it can catch. “Bigger is better” => collecting area What is the purpose of a telescope? 2. In addition to gathering light, a telescope allows a more detailed view of the structure of a celestial object and/or to discern the presence of multiple objects. This is called the telescope’s ANGULAR RESOLUTION A Telescope’s Diffraction Limit The ANGULAR RESOLUTION of a (single) telescope is always limited by its DIFFRACTION LIMIT. Minimum angular separation of two objects that can be seen as separate In radians Remember there are 206265 seconds of arc in one radian (a useful number to remember). Diffraction Limit In radians Example: Palomar 5m telescope The diameter of the telescope is 5 m = 500 cm Let’s find the diffraction limit at 500 nm. 1.22 X 500 nm X 10-7 cm/nm = 1.22 x 10-7 radians Θ= 500 cm = 1.22 x 10-7 radians X 206265 arcsec/radian = 0.025 arc seconds A ground-based optical telescope is limited by atmospheric turbulence => “atmospheric seeing” The “seeing” of an image The “seeing” of an image is a measure of its quality or sharpness. The seeing is always bigger than either (1) the diffraction limit or (2) the atmospheric seeing, whichever is greater. The “seeing” of an image • The seeing of an image is a measure of its quality or sharpness. • Because stars are so far away, they appear as points of light in our images. • The seeing then is the angular extent of a star in an image. • The seeing is always bigger than either (1) the diffraction limit or (2) the atmospheric seeing, whichever is greater. The Palomar 5-meter Hale Telescope • • • • Located in northern San Diego County, California Owned and operated by Caltech with various partners 5m telescope “first light” in 1949 Useful especially for spectroscopy, infrared imaging and adaptive optics/high resolution studies. The Palomar 5-meter Hale Telescope The Palomar 5-meter Hale Telescope High-Resolution Astronomy Solutions: • Put telescopes on mountaintops, especially in deserts • Put telescopes in space • Active optics – control mirrors based on temperature and orientation Images and Detectors Image processing by computers can sharpen images: “Adaptive optics”: take images rapidly; correct for “twinkling” caused by turbulence in the Earth’s atmosphere as light waves travel through it. The PHARO Adaptive Optics System Large Optical/IR telescopes Telescope Location Diameter Access Hubble space 2.4 m National/international VLT Chile 4x8m Europe Keck Mauna Kea 2 x 10 m Caltech/U California/Hawaii Gemini Mauna Kea and Chile 2x8m National/international Subaru Mauna Kea 7m Japan, U Hawaii Magellan Chile 2 x 6.5 m Carnegie, Harvard, MIT, Michigan, Chile Palomar Calif. 5m Caltech, JPL, National Access to some telescopes is restricted to astronomers from certain countries/institutions Telescopes European Southern Observatory Kitt Peak Natl Observatory Gemini Observatory Large optical telescopes Telescope Location Best seeing Limit Hubble space ~0.1” Near diffraction limit Typical ground based optical Palomar or Arizona ~1” Atmospheric seeing Best ground based optical Mauna Kea or Chile ~0.7” Atmospheric seeing + “conventional” Mauna Kea or adaptive optics Chile ~0.4” Some correction for atmospheric seeing + best/future adaptive optics ~0.1” Maximum correction for atmospheric seeing Mauna Kea or Chile Diffraction Limit In radians Example: Hubble Space Telescope HST The diameter of the telescope is 2.4 m = 240 cm Let’s find the diffraction limit at 500 nm. 1.22 X 500 nm X 10-7 cm/nm = 2.54 x 10-7 radians Θ= 240 cm = 2.54 x 10-7 radians X 206,265 arcsec/radian = 0.05 arc seconds Remember there are 206,265 seconds of arc in one radian (a useful number to remember). Electromagnetic spectrum Observing the universe Optical light: • Light from stars • Bright lines from ionized (hot) gas near very hot stars and supermassive black holes in galactic nuclei We need other telescopes to reveal: cold gas, cool gas, superhot gas, dust, and non-thermal sources (ones which don’t follow Wien’s law) Geometry, distances and time: HW#2 http://www.astro.ucla.edu/~wright/CosmoCalc.html Electromagnetic spectrum Not all electromagnetic radiation reaches the surface of the Earth; some blocked by atmosphere => Space telescopes (Hubble, Chandra, Spitzer, Compton Gamma Ray Observatory, etc.) Near Far IR 1 mm 10cm 10m 10 microns => 100 microns to 1 mm = Far IR/Submillimeter Radio Astronomy R-M-S: Radio – millimeter – submillimeter wavelengths Radio: Meter to centimeter wavelength 1 mm 1 cm 1 meter • Long wavelengths (relative to IR/opt/UV/X-rays) • By Wien’s law, we expect cold temperatures (partly true) • But also, not all radiation is thermal (i.e. follows Wien’s law and reflects the object’s temperature) •Synchrotron radiation •Bremsstrahlung radiation The Arecibo radio telescope Part of the National Astronomy and Ionosphere Center The Arecibo Observatory Located in northwestern Puerto Rico in a region of karst limestone sinkholes and “haystacks”. “light” of centimeter wavelengths How Arecibo works Inside the radome Rays are reflected off two additional mirrors in the radome. Diffraction Limit Example: Arecibo 305 meter telescope The diameter of the telescope is 305 m = 30500 cm Let’s find the diffraction limit at 21 cm. 1.22 X 21 cm Θ= = 8.4 x 10-4 radians 30500 cm = 8.4 x 10-4 radians X 206,265 arcsec/radian = 173 arc seconds = 2.9 arc minutes This is huge compared to Hubble’s 0.1 arcsec! Observing the Moon with Arecibo • At a wavelength of 21 cm => 1400 MHz (1.4 GHz) => a single pointing exposure with Arecibo detects photons from within an area about 3.5 arcminutes in diameter • Each pointing contributes one “pixel” to the final image • To contruct a full image of the Moon requires many pixels Observing the Moon with Arecibo • At a wavelength of 21 cm => 1400 MHz (1.4 GHz) => a single pointing exposure with Arecibo detects photons from within an area about 3.5 arcminutes in diameter • The final image covers the Moon but is blurry compared on an optical image. • Arecibo can detect faint objects but it has limited Light bucket angular resolution. ALFA: Arecibo L-band Feed Array “Camera” • 7 pixels (beams) •With ALFA, we can map wide areas of the sky 7 times faster than with a single pixel. •Because Arecibo is so big, we can detect very faint signals. ALFALFA: The Arecibo Legacy Fast ALFA Survey • Each pixel subtends about 3.5 arcminutes; this means that most galaxies are “unresolved”: • We detect the galaxy’s signal but, in most cases, do not get an image (the source is “unresolved”). • For each galaxy, we record its position, flux (intensity of signal) and the frequency at which the signal appears = Doppler shift. • The radio signal detected by ALFALFA arises from cool hydrogen gas in the galaxy’s interstellar medium. • Most sources detected by ALFALFA are clearly identified with a galaxy which also is seen at optical wavelengths (starlight). • But, some ALFALFA sources have no optical counterpart => “starless” galaxies… And maybe “totally dark” galaxies Radio Astronomers need Big Telescopes In radians How big would a radio telescope have to be to have a diffraction limit of 1 arc second at a wavelength of 21 cm? Θ = 1.22 X 21 cm Diameter cm = 1 arcsec/206,265 arcsec/radian Diam (cm) = 1.22 X 21 X 206,265 = 5.2 X 106 cm = 52 km (!) How can we possibly build a telescope that big???! Aperture Synthesis or Interferometry Sir Martin Ryle: 1974 Nobel prize in physics Use an array of smaller telescopes to achieve the image detail of a larger one that covers (sparsely) the area of the array. Interferometry Combine information from several widely-spread radio telescopes as if they came from a single dish • Resolution will be that of dish whose diameter = largest separation between dishes (“aperture synthesis”) Aperture Synthesis or Interferometry Sir Martin Ryle: 1974 Nobel prize in physics Use an array of smaller telescopes to achieve the image detail of a larger one that covers (sparsely) the area of the array. Aperture Synthesis or Interferometry Very Long Baseline Array Very Large Array Part of the National Radio Astronomy Observatory The Karl Jansky Very Large Array • 27 antennas, each one 25 m (85 ft) in diameter • Array in “Wye” (Y) shape; 4 configurations of “Wye” from compact to very spread out. • Located 70 miles west of Socorro New Mexico, which is about 70 miles south of Albuquerque. • Part of the National Radio Astronomy Observatory. The Karl Jansky Very Large Array • The resolution of the VLA is set by the size of the array. • 4 “configurations”, from compact to spread out by up to 22 miles. • Detects radio waves which vary from 0.7 cm wavelength to 4 meters wavelength achieving resolution of 0.05 arcsec (shortest wavelength) to 700 arcsec (longest wavelength) Very Long Baseline Array • The VLBA is the world’s only dedicated VLBI array. – Full complement of instrumentation – Time critical images of motions and source evolution – Unparalleled astrometry (microarcsec accuracy) 10 25-meter antennas, from Hawaii to St. Croix Angular resolution: 0.17 milliarcsec (0.00017 arcsec; shortest wavelength) to 22 milliarcsec (= 0.022 arcsec; longest wavelength) ALMA: Under construction The real thing…. Artist’s conception Atacama Large Millimeter Array • • • • • • Partners: North America, Europe, East Asia > 50 antennas, each 12-m in diameter Operational at millimeter to submillimeter wavelengths Located at 5000 m elevation in Atacama altiplano Scheduled for completion 2014 with ~60 antennas Partial array science around end of 2011 (starting now!) Will have angular resolution comparable to Hubble (0.1 arcsec) IRAS (1980s) Infrared Chandra (operational) X-rays Spitzer (partially operational) Infrared Hubble (operational;launched ) Optical PE #3: Multiwavelength telescopes Name Fermi Chandra GALEX HST Spitzer Herschel WISE ALMA EVLA Arecibo Wavelength range Diameter Location Main science Telescopes across the E-M spectrum Name Wavelength range Diameter Location Main science Gamma ray (complex) Low earth Time domain Chandra X-ray (complex) Elliptical orbit Imaging/spect GALEX 125-280 nm 0.5m Low earth Imaging/spect HST UV/opt/NIR 2.4m Low earth Imaging/spect NIR/MIR 0.9m Earth trailing Imaging/spect Herschel 60-670 μm 3.5m L2 (Lagrange point) Imaging/spect WISE 3.4-22 μm 0.4m Low earth Imaging ALMA 350μm–10mm 54 x 12m 5000 m in Chile Continuum/spect EVLA 7mm to 1m 27 X 25m 2124 m in NM Continuum/spect Arecibo 2 cm to 1 m 305 m Puerto Rico Pulsars; HI; Solar system radar Fermi Spitzer Hubble Space Telescope • • • • 2.4 meter reflecting telescope Image resolution ~ 0.1 arc second Deployed in low Earth orbit on 25 April 1990 Servicing missions: replace instruments/gyros • SM04 completed 2009 • Current instruments • Advanced Camera for Surveys (ACS: optical - fixed by SM04) • Space Telescope Imaging Spectrograph (STIS: UV – fixed SM04) • Wide Field Camera 3 (WFC3: UV to optical) • Cosmic Origins Spectrograph (COS: UV) • Near Infrared Camera & Multiobject Spectrograph (NICMOS: infrared; currently not functioning because of power issues) Spitzer Space Telescope • Launched on 25 August 2003 • 0.85-m (85 cm) telescope (largest ever in space!) • Observations at wavelength range from 3 to 180 microns • Because infrared radiation arises principally from thermal sources (“heat”), the telescope itself must be kept very cold. • Drifting heliocentric orbit (“deep space”) with solar shield => cannot be serviced (~5 year lifetime). • Ran out of liquid helium cryogen in 2009; now in “warm mission”, observing only in the 3-5 micron wavelength range Chandra X-ray Observatory • 1.2 meter X-ray telescope • Normal mirrors would absorb X-ray photons, so special design required • Wavelength range: 0.124 to 12.4 nanometers • Deployed in highly elliptical Earth orbit (reaches 1/3 of way to Moon) on 23 July 1999 • Cannot be serviced • Originally designed for min. 5 year mission, but still working very well. •Current instruments •Advanced CCD Imaging Spectrometer (ACIS) •High Resolution Camera (HRC) Telescope characteristics • • • • • • • • Aperture size (collecting area, diffraction limit) Wavelength/frequency coverage Elevation/transparency of atmosphere Angular resolution/point spread function Field of view Spectral bandwidth Spectral resolution Sampling rate (time domain) PE#3: Sextans A Radiation from galaxies In the optical regime, we detect the integrated starlight. I Thermal emission = black body radiation 3 2h 1 I()= c2 exp(h/kT) - 1 Typical spectrum of active galaxy, i.e. one with accreting supermassive black hole in its nucleus Bolometric magnitudes • Optical astronomers use λ; radio astronomers use . In fact: Apparent brightness • In the case of stars, we often use the bolometric magnitude, which is the magnitude integrated over all wavelengths. mbol = mV + B.C. Where the bolometric correction is defined as a function of Teff. B.C ~ 0 for F to G stars => peak in the V-band. Interstellar Dust • Probably formed in the atmospheres of cool stars • Mostly observable through infrared emission - very cold < 100 K. • Radiates lots of energy - surface area of many small dust particles adds of to very large radiating area • Infrared and radio emissions from molecules and dust are efficiently cooling gas in molecular clouds. • Whispy nature indicates turbulence in ISM IRAS (infrared) image of infrared cirrus of interstellar dust. Sampling the SED U B V R I 1 nm = 10 Å Adding the NIR PE#3: Barnard 68 B V Z K H J The color index Suppose we observe a star in two standard wavelength ranges centered at 1 and 2. The “color index” is proportional to the ratio of their observed fluxes (brightnesses). m2 – m1 = 2.5 log [F(1)/ F(2)] flux color index λ→ The color index is a “cheap” measure of spectral class; Careful design of filters to give max. info.