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

+ d2(,)/dt2 = 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 1906

100-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 “fl” of

the telescope.

= 206265”/fl(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

sin11 = sin22

2/1 = 12

sin2 = sin1 /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

10-20 m UV?

Ground

LSST, GSMT (GMT,TMT,EELT….)

TMT

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 / cm2 / 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 +/- 0.67 68% “ “ “ +/- 1.00 90% “ “ “ +/- 1.69 95 % “ “ “ +/- 1.96 99 % “ “ “ +/- 2.58 99.6% “ “ “ +/- 3.00

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/cm2/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 cm2

t = integration time in sec

W = bandwidth in A

O = Object rate (cts/s/cm2/A)

B = Sky (background) rate

D = dark rate

R = read noise

S/N = OAtW/((O+B)AtW + Dt + R2)1/2

Special Cases

Background limited (B >> D or R)

S/N = O/(O+S)1/2 x (AtW)1/2

Detector limited (R2 >> 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 = 1.000

Divide science images by FF, filter by filter.

Apply other routines as necessary.

Astronomical PhotometryFor 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 log10(vcts/sec) + constant

V = v + C1(B-V) + kVx + C2 ……

x = sec(zenith distance) = airmass

(B-V) = C3(b-v) + C4 + kBVx + ….