Download Astro 3303 - Cornell Astronomy

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.