1. The Imaging Chain in Optical Astronomy

2. Review and Overview “Imaging Chain” includes these elements: 1.energy source 2.object 3.collector 4.detector (or sensor) 5.processor 6.display 7.analysis 8.storage (if any)

3. Optical Imaging Chain 1: source 2: object 3: collector 4: sensor 5: processing 6: display 7: analysis

4. Source and/or Object In astronomy, the source of energy (1) and the object (2) are almost always one and the same! i.e., The object emits the light –Examples: Galaxies Stars –Exceptions: Planets and the moon Dust and gas that reflects or absorbs starlight

5. Optical Imaging Chain in Astronomy 1: source 2: object 3: collector 4: sensor 5: processing 6: display 7: analysis 8: storage or

6. Optical Imaging Chain in Radio Astronomy 1,2 3,4 5 6,7

7. Specific Requirements for Astronomical Imaging Systems Requirements always conflict –Always want more than you can have  must “trade off” desirable attributes  Deciding the relative merits is a difficult task  “general-purpose” instruments (cameras) may not be sufficient Want simultaneously to have: –excellent angular resolution AND wide field of view –high sensitivity AND wide dynamic range Dynamic range is the ability to image “bright” and “faint” sources –broad wavelength coverage AND ability to measure narrow spectral lines

8. Angular Resolution vs. Field of View Angular Resolution: ability to distinguish sources that are separated by small angles –Limited by: Optical Diffraction Sensor Resolution Field of View: angular size of the image field –Limited by: Optics Sensor Size (area)

9. Sensitivity vs. Dynamic Range Sensitivity –ability to measure faint brightnesses Dynamic Range –ability to image “bright” and “faint” sources in same system

10. Wavelength Coverage vs. Spectral Resolution Wavelength Coverage –Ability to image over a wide range of wavelengths –Limited by: Spectral Transmission of Optics (Glass cuts off UV, far IR) Spectral Resolution –Ability to detect and measure narrow spectral lines –Limited by: “Spectrometer” Resolution (number of lines in diffraction grating)

11. Optical Collector (Link #3)

12. Optical Collection (Link #3): Refracting Telescopes Lenses collect light BIG disadvantages –Chromatic Aberrations (due to dispersion of glass) –Lenses are HEAVY and supported only on periphery Limits the Lens Diameter Largest is 40" at Yerkes Observatory, Wisconsin http://astro.uchicago.edu/vtour/40inch/kyle3.jpg

13. Optical Collection (Link #3): Reflecting Telescopes Mirrors collect light Chromatic Aberrations eliminated Fabrication techniques continue to improve Mirrors may be supported from behind  Mirrors may be made much larger than refractive lenses

14. Optical Reflecting Telescopes Concave parabolic primary mirror to collect light from source –modern mirrors for large telescopes are thin, lightweight & deformable, to optimize image quality 3.5 meter WIYN telescope mirror, Kitt Peak, Arizona

15. Thin and Light (Weight) Mirrors Light weight  Easier to point –“light-duty” mechanical systems  cheaper Thin Glass  Less “Thermal Mass” –Reaches Equilibrium (“cools down” to ambient temperature) quicker

16. Hale 200" Telescope Palomar Mountain, CA http://www.astro.caltech.edu/observatories/palomar/overview.html http://www.cmog.org/page.cfm?page=374

17. 200" mirror (5 meters) for Hale Telescope Monolith (one piece) Several feet thick 10 months to cool 7.5 years to grind Mirror weighs 20 tons Telescope weighs 400 tons “Equatorial” Mount –follows sky with one motion

18. Keck telescopes, Mauna Kea, HI http://www2.keck.hawaii.edu/geninfo/about.html

19. 400" mirror (10 meters) for Keck Telescope 36 segments 3" thick Each segment weighs 400 kg (880 pounds) –Total weight of mirror is 14,400 kg (< 15 tons) Telescope weighs 270 tons “Alt-azimuth” mount (left-right, up-down motion) –follows sky with two motions + rotation

20. Basic Designs of Optical Reflecting Telescopes 1.Prime focus: light focused by primary mirror alone 2.Newtonian: use flat, diagonal secondary mirror to deflect light out side of tube 3.Cassegrain: use convex secondary mirror to reflect light back through hole in primary 4.Nasmyth (or Coudé) focus (coudé  French for “bend” or “elbow”): uses a tertiary mirror to redirect light to external instruments (e.g., a spectrograph)

21. Prime Focus f Sensor Mirror diameter must be large to ensure that obstruction is not significant

22. Newtonian Reflector Sensor

23. Cassegrain Telescope Sensor Secondary Convex Mirror

24. Feature of Cassegrain Telescope Long Focal Length in Short Tube Location of Equivalent Thin Lens f

25. Coudé or Nasmyth Telescope Sensor

26. Optical Reflecting Telescopes Schematic of 10-meter Keck telescope (segmented mirror)

27. Large Optical Telescopes Telescopes with largest diameters (in use or under construction : –10-meter Keck (Mauna Kea, Hawaii) –8-meter Subaru (Mauna Kea) –8-meter Gemini (twin telescopes: Mauna Kea & Cerro Pachon, Chile) –6.5-meter Mt. Hopkins (Arizona) –5-meter Mt. Palomar (California) –4-meter NOAO (Kitt Peak, AZ & Cerro Tololo, Chile) http://seds.lpl.arizona.edu/billa/bigeyes.html Summit of Mauna Kea, with Maui in background Keck telescope mirror (note person)

28. Why Build Large Telescopes? 1.Larger Aperture  Gathers MORE Light –Light-Gathering Power  Area –Area of Circular Aperture =  D 2 / 4  D 2 D = diameter of primary collecting element 2.Larger aperture  better angular resolution –recall that:

29. Why Build Small Telescopes? 1.Smaller aperture  collects less light  less chance of saturation (“overexposure”) on bright sources 2.Smaller aperture  larger field of view (generally) –Determined by “F ratio” or “F#” f = focal length of collecting element D = diameter of aperture

30. F Ratio: F# F# describes the ability of the optic to “deflect” or “focus” light –Smaller F#  optic “deflects” light more than system with larger F# Small F# Large F#

31. F# of Large Telescopes Hale 200" on Palomar: f/3.3 –focal length of primary mirror is: 3.3  200" = 660" = 55'  16.8 m –Dome must be large enough to enclose Keck 10-m on Mauna Kea: f/1.75 –focal length of primary mirror is: 1.75  10m = 17.5 m  58 m

32. F Ratio: F# Two reflecting telescopes with different F# and same detector have different “Fields of View”: Small F#Large F# large  small 

33. Sensors (Link #4)

34. Astronomical Cameras Usually Include: 1.Spectral Filters –most experiments require specific wavelength range(s) –broad-band or narrow-band 2.“Reimaging” Optics –enlarge or reduce image formed by primary collecting element 3.Light-Sensitive Detector: Sensor

35. Astronomical Sensors Most common detectors: –Human Eye –Photographic Emulsion film plates –Electronic Sensors CCDs

36. Angular Resolution Fundamental Limit due to Diffraction in “Optical Collector” (Link #3) But Also Limited by Resolution of Sensor!

37. Charge-Coupled Devices (CCDs) Standard light detection medium for BOTH professional and amateur astronomical imaging systems –Significant decrease in price numerous advantages over film: –high quantum efficiency (QE) meaning most of the photons incident on CCD are “counted” –linear response measured signal is proportional to number of photons collected –fast processing turnaround (CCD readout speeds ~1 sec) NO development of emulsion! –regular grid of sensor elements (pixels) as opposed to random distribution of AgX grains –image delivered in computer-ready form

38. CCD Basics Light-sensitive electronic element based on crystalline silicon –crystal = “lattice” of atoms at regular spacings –acts as though electrons have two states: “bound” to atom “free” to roam through lattice

39. CCD Basics Incident photon adds energy to electron to “kick” it up into the “free” states –energy of photon must be sufficiently large for electron to “reach” the free states –to be absorbed, the photon wavelength must be less than maximum max  1100 nm (near infrared) Electrons in “Free” States (“conduction band”) Electrons in “Bound” States (“valance band”) Energy photon

40. CCD Basics Silicon structure is divided into pixels –e - transferred and “counted” one pixel at a time http://www.byte.com/art/9510/img/505099d2.htm

41. Sensor Resolution Obvious for Electronic Sensors (e.g., CCDs) xx Elements have finite size Light is summed over area of sensor element (“integrated”) Light from two stars that falls on same element is added together stars cannot be distinguished in image!

42. Same Effect in Photographic Emulsions More difficult to quantify Light-sensitive “grains” of silver halide in the emulsion Placed “randomly” in emulsion “Random” sizes “large” grains are more sensitive (respond to few photons) “small” grains produce better resolution

43. Photographic techniques: silver halide Film –Emulsion on “flexible” substrate –Still used by amateurs using sensitive film B&W and color Special treatment to increase sensitivity Photographic Plates –Emulsion on glass plates –Most common detector from earliest development of AgX techniques until CCDs in late 70’s

44. Eye as Astronomical Detector Eye includes its own lens –focuses light on retina ( “sensor”) When used with a telescope, must add yet another lens –redirect rays from primary optic –make them parallel (“collimated”) rays appear to come from “infinity” (infinite distance away) –reimaging is performed by “eyepiece”

45. Eye with Telescope Without Eyepiece With Eyepiece Light entering eye is “collimated”

46. Eye as Astronomical Detector Point sources (stars) appear brighter to eye through telescope Factor is –D is telescope diameter –P is diameter of eye pupil –Magnification should make light fill the eye pupil (“exit pupil”) Extended sources (for example, nebulae) do not appear brighter through a telescope –Gain in light gathering power exactly compensated by image magnification, spreads light out over larger angle.

47. Atmospheric Effects on Image Large role in ground-based optical astronomy –scintillation modifies source angular size twinkling of stars = “smearing” of point sources –extinction reduces light intensity atmosphere scatters a small amount of light, especially at short (bluer) wavelengths water vapor blocks specific wavelengths, especially near-IR –scattered light produces interfering “background” astronomical images are never limited to light from source alone; always include “source” + “background sky” “light pollution” worsens sky background

48. Scattering “Wavelength Dependent” –Depends on color of light –Long wavelengths are scattered “less”

49. Scattering by Molecules Molecules are SMALL “Blue” light is scattered MUCH more than red light –Reason for BOTH blue sky (blue light scattered from sun in all directions) red sunset (blue light is scattered out of the sun’s direct rays)

50. Scattering by Dust Dust particles are MUCH larger than molecules –e.g., from volcanos, dust storms Blue light is scattered by dust “somewhat more” than red light

51. Link #5: Image Processing

52. Formerly: performed in darkroom –e.g., David Malin’s “Unsharp Masking” Subtract a blurred copy from a “sharp” positive (or, add a blurred negative to a “sharp” positive) Now performed in computers, e.g., –contrast enhancement –“sharpening” –“normalization” (background division) –…

53. Example of Unsharp Masking http://www.hawastsoc.org/messier/fslide53.html UnprocessedAfter Unsharp Masking http://www.seds.org/messier/m/m042.html n.b., Increased visibility of fine structure in bright and dark regions of “cloud” after unsharp masking

54. Blurring vs. Sharpening Blurring: –Local “Averaging” of Pixels in Scene –“Averages out” fine detail in image more than large-scale structure Sharpening: –“Inverse” of Blurring  Local “Differencing” of Pixels

55. Image Processing to Correct for: Atmosphere (to extent possible) –e.g., images obtained of object at different “heights” in sky exhibit different atmospheric “extinction” –images usually can be corrected to compare brightnesses CCD defects and artifacts –“dark current” Pixel gives output response even when not exposed to light –Bad pixels Due to manufacturing flaws “Dead”, “Hot”, “Flickering” (time-variable response) –Variations in pixel-to-pixel sensitivity every pixel has its own Quantum Efficiency (QE) Characterized by measuring response to uniform “flat field” and subsequently “divided out”

56. Links #6 and #7 Image Display and Analysis

57. Image Display and Analysis This step often is where astronomy really begins. Type and extent of display and analysis depends on purpose of imaging experiment Common examples: –evaluating whether an object has been detected or not –determining total CCD signal (counts) for an object, such as a star –determining relative intensities of an object from images at two different wavelengths –determining relative sizes of an extended object from images at two different wavelengths

58. Link #8: Storage

59. Storage Glass plates –Requires MUCH climate-controlled storage space –Expensive to store and retrieve –available to one user at a time –now being “digitized” (scanned), as in the archive you use with DS9 Digital Images –Lots of disk space –cheaper all the time –available to many users

60. Objects with Large Luminosities can be Detected at Large Distances Lines of constant apparent brightness

61. Distant Objects are Usually Physically Larger Lines of constant angular size

62. Angular Sizes Span Wide Range