Telescopes and Astronomical Observations Ay16 Lecture 5 Feb 14, 2008

Outline: What can we observe? Telescopes Optical, IR, Radio, High Energy ++ Limitations Angular resolution Spectroscopy Data Handling

A telescope is an instrument designed for the observation of remote objects and the collection of electromagnetic radiation. "Telescope" (from the Greek tele = 'far' and skopein = 'to look or see'; teleskopos = 'far-seeing') was a name invented in 1611 by Prince Frederick Sesi while watching a presentation of Galileo Galilei's instrument for viewing distant objects. "Telescope" can refer to a whole range of instruments operating in most regions of the electromagnetic spectrum.

Telescopes are “Tools” By themselves, most telescopes are not scientfically useful. They need yet other tools a.k.a. instruments.

What Can We Observe? Brightness (M) + dM/dt = Light Curves, Variability + dM/d = Spectrum or SED + dM/d /dt = Spectral Variability Position + d( ,  )/dt = Proper Motion + d 2 ( ,  )/dt 2 = Acceleration Polarization

“Instruments” Flux detectors Photometers / Receivers Imagers Cameras, array detectors Spectrographs + Spectrometers “Spectrophotometer”

Aberrations Spherical Coma Chromatic Field Curvature Astigmatism

Mt. Wilson & G. E. Hale 60-inch inch 1917

Edwin Hubble at the Palomar Schmidt Telescope circa 1950

Telescope Mirrors Multiple designs Solid Honeycomb Meniscus Segmented

Focal Plane Scale Scale is simply determined by the effective focal length “f l ” of the telescope. = ”/f l (mm) arcsec/mm * Focal ratio is the ratio of the focal legnth to the diameter

Angular Resolution The resolving power of a telescope (or any optical system) depends on its size and on the wavelength at which you are working. The Rayleigh criterion is sin (  ) = 1.22 /D where  is the angular resolution in Radians

Airy Diffraction Pattern * more complicated as more optics get added…

Encircled Energy Another way to look at this is to calculate how much energy is lost outside an aperture. For a typical telescope diameter D with a secondary mirror of diameter d, the excluded energy is x( r) ~ [5 r (1- d/D)] -1 where r is in units of /D radians  a 20 inch telescope collects 99% of the light in 14 arcseconds

2 Micron All- Sky Survey 3 Channel Camera

Silicon Arrays --- CCDs

CCD Operation Bucket Brigade

FAST Spectrograph

Simple Fiber fed Spectrograph

Hectospec (MMT )

Holmdel Horn

GBT

Astronomical Telescopes & Observations, continued Lecture 6 The Atmosphere Space Telescopes Telescopes of the Future Astronomical Data Reduction I.

Atmospheric transparency

Hubble

Ground vs Space

Adaptive Optics

Chandra X-Ray Obs

Grazing Incidence X-ray Optics Total External Reflection

X-Ray Reflection Snell’s Law sin  1  1 = sin  2  2  2 /  1 =  12 sin  2 = sin  1 /  12 Critical angle = sin  C =  12 --> total external reflection, not refraction

GLAST A Compton telecope

Compton Scattering

LAT GBM

The Future? Space JWST, Constellation X m UV? Ground LSST, GSMT (GMT,TMT,EELT….)

TMT

GMT

EELT = OWL

OWL Optical Design

JWST

ConX

Chinese Antarctic Astronomy

Astronomical Data Two Concepts: 1. Signal-to-Noise 2. Noise Sources

Photon Counting Signal O = photons from the astronomical object. Usually time dependent. e.g. Consider a star observed with a telescope on a single element detector O = photon rate / cm 2 / s / A x Area x integration time x bandwidth = # of photons detected from source

Noise N = unwanted contributions to counts. From multiple sources (1) Poisson(shot) noise = sqrt(O) from Poisson probability distribution (Assignment: look up Normal = Gaussan and Poisson distributions)

Poisson Distribution

Normal=Gaussian Distribution The Bell Curve

Normal = Gaussian 50% of the area is inside +/  68% “ “ “ +/  90% “ “ “ +/  95 % “ “ “ +/  99 % “ “ “ +/  99.6% “ “ “ +/  of the mean

(2) Background noise from sky + telescope and possibly other sources Sky noise is usually calculated from the sky brightness per unit area (square arcseconds) also depends on telescope area, integration time and bandpass B = Sky counts/solid angle/cm 2 /s/A x sky area x area x int time x bandwidth

Detector Noise (3) Dark counts = D counts/second/pixel (time dependent) (4) Read noise = R (once per integration so not time dependent)

So if A = area of telescope in cm 2 t = integration time in sec W = bandwidth in A O = Object rate (cts/s/cm 2 /A) B = Sky (background) rate D = dark rate R = read noise S/N = OAtW/((O+B)AtW + Dt + R 2 ) 1/2

Special Cases Background limited (B >> D or R) S/N = O/(O+S) 1/2 x (AtW) 1/2 Detector limited (R 2 >> D or OAtW or BAtW) S/N = OAtW/R (e.g. high resolution spectroscopy)

CCD Data Image data cts/pixel from object, dark, “bias” Image Calibration Data bias frames flat fields dark frames (often ignored if detector good)

Image Display Software SAODS9 Format.fits

NGC1700 from Keck

Spectra with LRIS on Keck

Bias Frame gives the DC level of the readout amplifier, also gives the read noise estimate.

Flat Field Image through filter on either twilight sky or dome

Image Reduction Steps Combine (average) bias frames Subtract Bias from all science images Combine (average) flat field frames filter by filter, fit smoothed 2-D polynomial, and divide through so average = Divide science images by FF, filter by filter. Apply other routines as necessary.

Astronomical Photometry For example, for photometry you will want to calibrate each filter (if it was photometric --- no clouds or fog) by doing aperture photometry of standard stars to get the cts/sec for a given flux Then apply that to aperture photometry of your unknown stars. NB. There are often color terms and atmospheric extinction.

Photometry, con’t v = -2.5 x log 10 (v cts/sec ) + constant V = v + C 1 (B-V) + k V x + C 2 …… x = sec(zenith distance) = airmass (B-V) = C 3 (b-v) + C 4 + k BV x + ….