observations & instrumentation ii: spectroscopy

8/12/2019 Observations & Instrumentation II: Spectroscopy

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ASTR 3520

Observations & I nstrumentation I I :Spectroscopy

Lecture 1

Introduction


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OverviewJohn Bally C323A Duane 492 5786

[email protected]

[email protected]

Office hours: Th after class (2:00 PM)

Wed (2:00 PM)

Adam GinsburgC329 Duane 303 667 3805

[email protected]

Office Hours: Mon, Tues 11:00 AM

or by appointment

Student & Teacher Introductions:


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Organization Review course structure, content, and Syllabus

Observing Projects: Stellar, nebular spectroscopy,semester projects, labs, homework.

Apache Point Observatory Field Trip:

- 5 - 6 days/ 4 - 5 nights- Covered by Course Fees

- VLA, NSO, APO

- Last week of Oct. (depends on TAC)

Observing Proposals for Semester project due end of Sept.

24 Observing Groups 5 groups / 3 to 4 each.

- Each group must have at last 1 experienced observer

Start spectrograph overview (once-over lightly)


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Spectroscopy: Astronomy => Astrophysics

Light as a wave phenomenon: = cGeometrical optics => wave opticsDiffraction~ / DInterference:

n= D sin n = 1,2,3, Deep insights into the nature of atoms, molecules:

Discrete wavelengths => Discrete energy levels

Electrons stable only in certain orbits.

Interference of electron waves!

= h / p = h / mv :de Broglie wavesAll matter has wave-like behavior on sufficiently

small scale!


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Telescope

Focal Plane

Slit

Spectrograph

Spectrograph

collimator

Dispersing element

camera

detector


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SBO Spectrograph overview

Slit & Decker:Restrict incoming light

Spatial direction vs. Spectral direction

Collimator & Camera:

Transfer image of slit onto detector.

Grating:

Disperse light: dispersion => spectral resolution

What determines spectral resolution & coverage?

- Slit-width- Grating properties: Ngroves , order number

- Camera / collimator magnification (focal length ratio)

- Detector pixel size and number of pixels.


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Types of Spectroscopy Electromagnetic Waves: Emission, absorption

Visual, near-IR., FIR, Radio, UV/X-ray, gamma-ray

- Solids, liquids, gasses, plasmas- Emission, absorption

- Spectral line, molecular bands, continua:

- Thermal (~LTE, blackbody, grey-body):

- Non-thermal (masers, synchrotron, )- Electronic, vibrational, rotational transitions.

- Effects of B (Zeeman), E ( Stark), motion (Doppler),

pressure (collisions), natural life-time (line widths)

- Radiative Transfer (optical depth)

Other types (not covered in this course):

NMR

Raman

Phosprescence / Fluorecence

Astro-particle


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Review of Some Basics c = x Angular resolution: = 1.22 / D radians

206,265 in a radian E = h F = L / 4 pd2 AZ, El, RA, Dec, Ecliptic, Galactic

Siderial time, Hour Angle

G = 6.67 x 10-8(c.g.s)

c = 3 x 1010 cm/sec,

k = 1.38 x 10-16

h = 6.626 x 10-27

mH~ mproton= 1.67 x 10-24gramsme= 0.91 x 10

-27grams

eV = 1.602 x 10-12erg

Luminosity of Sun = 4 x 1033erg/sec

Mass of the Sun= 2 x 10

33

grams


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The Physics of EM Radiation

Light: , - = c = 2.998 x 1010cm/s (in vacuum)- E = h Photon energy (erg)

1 erg sec-1= 10-7Watt

h = 6.626 x 10-27 (c.g.s)

1 eV = 1.602 x 10-12erg

- p = E / c = h / Photon momentum- = h / p = h / mv deBroglie wavelength

Planck Function: B(T) Emission, absorption, continua

Discrete energy levels: Hydrogen


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Refraction:

Snells Law: n1sin(d1) = n2sin(d2)

d2

d1 n1

n2

n1= refractive index in region 1

n2= refractive index in region 2

n = c / v = vacuum/medium


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Basic Lens formulae:


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Basic Mirror formulae:


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Diffraction:

Light spreads as = / dIn the `far field given by L = d2/

d

L


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2 slit interference

Constructive Destructive


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Anti-reflection coating

2 slit interference


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Multi-layer interference filter:


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Diffraction grating:


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Fermats Principle: d(optical path length) = 0

Diffraction grating:

order #

wavelength

groove spacing incidence anglediffraction angle


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CCD Imaging Review

Review CCD basics- How CCDs work

- CCD properties

Dark, flat, and bias frames Image-scales

- focal length, pixel-scale, FOV

Review photometry basics- The magnitude system

- Calibration

- Atmospheric effects; Air mass, color terms


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Subaru 8m (Mauna Kea): Suprime Prime Focus CCD Mosaic8192 x 8192 pixels using SITe chips (15 mm pixels)


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Typical

Raw image

With a CCD

Cosmic rays

Bad pixels

stars


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CCDs (Charge-Coupled Device)

Properties

- Quantum efficiency (QE):

=> 90%

- Gain:

G = e- /ADU

- Dark current:1 e-/ hr to 103e-/sec

thermal emission: => Cool to20 to150 C

- Read Noise:

amplifier read-out uncertainty3 e-to 100 e- per read

- Spatial uniformity:

Bad pixels, columns: ~


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CCDs

Properties

- Cosmic Rays:

5 to > 103 e-produced by each charged particle

usually effects 1 or few pixels.

non-gaussian charge distribution

(different from stellar image or PSF)- Well depth:

5 x 104to 106 e-

- Pixel size:

6 mm to 30 mm- Array size:512 x 512 to 4096 x 4096


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Dark current:

=> cooling


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MOSAIC CCD

On KPNO 0.9m

Vacuum Dewar

LN2(77K)

Controller

Filters & slider

V


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Charge Transfer0

510

05

10

510

0

V


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Charge Coupled Devices (CCDs)

Output amplifier


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Output amplifier


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Read



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Read



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CCD Corrections/Calibrations

Read noise: bias frames

- 0 second exposure

Dark frames:

- Same duration as science exposure with shutter closed

Flat fields:

- Dome flats- Twilight flats

- Super-sky flats

Standard stars

- At several air-masses

A = sec (z) = 1 / cos(z)

z


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Types of image combinations:

IRAF task: imarith image1 (+,-,*,/) image2 output

imcombine @list_in output

- Average: 1/N SI(n)- Mode: Most common data value

- Median: Value in middle of rangegood for rejection of outliers (e.g CRs)

Combine (median) 3,5,7,.. An odd #

- bias frames- flat frames


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Reduction:

I(raw) - median(bias)

I(reduced) =

norm [median(Flat bias)]

Note: Bias can be a Dark if hot pixels /or dark current is

large


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Flat Field Examplestar

star

cosmic ray

star

star

cosmic ray

Bias ordark levelRaw science frame

Dark subtracted frame

Hot pixels

i


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Flat Field Examplestar

star

cosmic ray

cosmic ray

Flat frame

Fl Fi ld E l


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Flat Field Examplecosmic ray

Flat frame

Normalized, dark subtracted, median of > 3 flat frames

1

Fl t Fi ld E lcosmic ray


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Flat Field Example

Normalized flat frame

1

star

cosmic ray

Science frame

star star

Reduced science frame

Ph B i


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Photometry Basics:

Vega magnitudes:

m() = -2.5 log [F() / FVega()]F() = Counts on source

FVega() = Counts on Vega

A = sec (z) = 1 / cos(z)z

T f S t


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Type of Spectra

Continuum:- Blackbody: B(T)- free-free, free-bound

- Non-thermal: Synchrotron radiation

- Compton scattering

Line & Band

E dipole, B diplole, E quadrupole

fine structure, hyperfine structure- electronic transitions

- vibrational transitions

- rotational transition


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Types of Spectra:

Nebulae

Stars

Hot,

Opaque

media


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The Planck Function: Black-body radiation

Rayleigh-Jeans:

Wien:

B(,T) = (2 p h3/ c2) e-h/kT

B(,T) = 2kT/2

(erg s-1cm-2Hz-1 2 psr-1)



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Wien Rayleigh-Jeans


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Spectrum of Hydrogen (& H-like ions)


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Lyman

Balmer

a

b

a

b


Ionization (n to infinity):

E = 13.6 eV

Transitions:

E = h= EuEl

Ionization at

E = 13.6 eV or less than = 912 Angstroms

= R [ 1/nl21 /nu

2]

R = 3.288 x 1015Hz


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Bohr model:

Allowed orbits

mvr = nh /2pCoulomb Force:

Ze2/ r2= mv2/r

Thus, (eliminate v)

r = Ze

2

/ mv

2

= n

2

h

2

/ 4 p2 Ze2mEnergy E = - (1/2) Ze2/ r = - 2 p2Z2e4m/ n2h2


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Adam Block: 16 Meade + SBIG ST10E + AO7


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The Orion Nebula (M42)


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O tli & G l


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Outline & Goals: Tues, 18 Sept

Summary of Kitt Peak Run &Heildelberg

Review Spectrum of Hydrogen

Spectroscopic `terms & terminology

(Ch 2, 3; HW #2)

Review Transitions (Ch 3): Einstein

A, B. Col l isional and radiativeexcitation

Spectral l ine formation & Radiative

Transfer basics


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Bohr model:

Allowed orbits

mvr = nh /2pCoulomb Force:

Ze2/ r2= mv2/r

Thus,

r = Ze2/ mv2 = n2h2/ 4 p2 Ze2mEnergy E = - (1/2) Ze2/ r = - 2 p2Z2e4m/ n2h2



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Lyman

Balmer

a

b

a

b

Sp y g (& )

Ionization (n to infinity):

E = 13.6 eV

Transitions:

E = h= EuEl

Ionization at

E = 13.6 eV or less than = 912 Angstroms

= R [ 1/nl21 /nu

2]

R = 3.288 x 1015Hz

Ionization cross-section or hydrogen


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Wavelength (1 / photon energy)

10-18

-3

Lyman lines

y g

13.6 eV = 912 Angstroms

Balmer lines

Atomic Structure


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Atomic Structure

Refinements to Bohr:Elliptical e-orbits

Integral of P in r and = lh l = 0,1,2, ,n-1Relativistic effects => l makes smallcorrection to E-levels

Space quantization: Orientation of orbits

m

Electron spin

Pauli: No 2 e- in same state.

Atomic Structure


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Atomic Structure

Atomic quantum numbers:n, l, m, s - completely specify state, E

n = 1, 2, 3, 4 .shell = K L M N .

max ne = 2 8 .

l = 0 1 2 3 4 .

s p d f g .

Selection rules:

Atomic Structure


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Atomic Structure Refinements to Bohr: n

Elliptical e-orbits: k

Space quantization: Orientation of orbits

w.r.t. magnetic field: m

Electron spin: s

Pauli: Ferminons:No 2 e- in same state: [n,k,m,s]

Shroedinger Wave function:

n => principle quantum number (radial)

l => orbital angular momentum 0, 1, nm => magnetic sublevels (degenerate if B=0)

s => electron spid +/- 1/2

Atomic Structure


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Atomic Structure

Multi-electron atoms/ions

Atomic quantum numbers: s = +/- 1/2n, l, m, s - completely specify state, E

l = 0, 1, , (n-1) m = 0, +/- 1, +/- 2, , +/- l

n = 1, 2, 3, 4 .

shell = K L M N .l = 0 0, 1 0,1,2 0,1,2,3

m = 0 0; -1,0,1 0;-2,-1,0,1,2 0;-3,-2,-1,0,1,2,3

max ne = 2 2+6 = 8 2+8+10 = 20 (s = +/- 1/2)

max l = 0 1 2 3 4

s s,p s,p,d s,p,d,e

Selection rules: Dl = +/- 1

n


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n

1

23

4


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Hydrogen Ha fine structure

Review Transitions (Ch 3): Einstein

A, B. Col l isional and radiative

excitation

Spectral l ine formation & Radiative

Transfer basics

Hydrogen energy levels showing allowed transitions


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Hydrogen energy levels showing allowed transitions

Hydrogen energy levels showing fine structure


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Hydrogen energy levels showing fine structure

l = 0 l = 1 l = 2

S P D

n=3

n=2

n=1

Ha

Lya

Fine structure const.:a= e2/ hc = 1/137Fine structure:dE / E ~ a4 ~ 5x10-5 eVSpin / orbit (l * s)

1s 2S1/2

2s2S1/2

3s 2S1/2

3p2Po3/2

3p2Po1/2

3d2D5/2

3d2D3/2

2p2Po3/2

2p2Po1/2

2s+1

J=L+S

Hyperfine structure:dE ~ 6x10-6 eVSelection Rules:Dl = 0, +/-1Dj = 0, +/-1even odd

L = [l(l+1)]1/2h/2pS = [s(s+1)]1/2h/2pJ = [J(J+1)]1/2h/2p

Einstein A & B coefficients: radiative processes


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Einstein A & B coefficients: radiative processes

Aul Blu BluBul

Aul

Bul

- Spontaneous decay

- Stimulated decay (prop to Flux)- Absrorption (prop to Flux)

u

l

Ionization & Excitation


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Ionization & Excitation

Radiation: Rr = sIrad Collisions: Rc = n s

Blu

Aul, Bul Cul

Clu

Rate Equations:

Vthermal~ (3 kT / 2 m)1/2

s= cross section


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Guidelines for 24 spectroscopy

Simple Models for Spectrum Formation

emission nebulaeabsorption line features

continuum processes

Theory of Spectrum Formationoptically thin and optically thick spectra

stellar spectra

Spectral Line Formation &

Radiative Transfer


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Emission Nebulae

Atoms in nebulae are excited by:

Incident photons

Collisions (high temperature or density) Excited atoms decay, emitting a photon of

the characteristic energy (a spectral line)

If the atoms are ionized, then the nebulawill emit free-bound radiation (i.e. Balmercontinuum) as well as spectral lines

Ionization cross-section or hydrogen


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Wavelength (1 / photon energy)

10-18

-3

Lyman lines

13.6 eV = 912 Angstroms

Emission Nebula


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Star emits continuum

optically thin nebula:

passes most wavelengths

- light at energy equal to an

atomic transition is absorbed

- that light is then reemitted in a

random direction (some of it

towards the observer)

- the nebula may be optically

thick at these wavelengths

Emission Nebula(photo-excited or photo-ionized)

The only light directed towardsthe observer is that which has

energy equal to the atomic

transitions in the nebula:

an emission spectrum

The Ring Nebula (M57): Planetary Nebula


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The Swan Nebula (M17): Emission Nebula


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M82Subaru 8-m (Mauna Kea)


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continuum

Emission line (Ha)

Absorption (dust, NaI, )

emission (stars)

M82UV (Galex)


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M8221 cm HI (VLA)


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M82

M81

NGC3077

M82 Ha


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M82 radio (6 cm)

M82 X ray


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M82 X-ray


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Absorption Features

Continuum light is emitted from a star (or

other source)

Intervening material absorbs light atwavelengths of atomic transitions, exciting

those atoms

Excited atoms reemit light, but in a randomdirection (not towards observer)

Absorption Feature:


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the observer sees all the

wavelengths except those

at the atomic transition energy

an absorption spectrum

- light at energy equal to an atomic

transition is absorbed

- that light is then reemitted in a

random direction

p

QSO Spectrum with IGM Absorption


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What Does an Absorption Spectrum Look Like in an Image?


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Quasar 3C273 Deneb

p p g

It looks almost identical to the background object!

All the absorption is in a few lines, the continuum

is relatively unchanged.

optically thick nebula:Dark or Reflection Nebula:


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p y

-observer cant see background

object (i.e. star) because light

has been scattered away

- dark nebula

dust scatters light

optically thick nebula:

-observer sees cloud shining in

scattered light (a continuum)

-reflection nebula

Continuum Phenomena:


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Reflection and Dark Nebulae

Dust scatters incident light

Not a line process, scatterscontinuum

reflection

dark

dark


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Basic Radiative Transfer Terms

MFP = mean free path (cm)

au= opacity (cm-1)cross section per unit volume (aka

absorptivity)

MFP = 1/au

tu = optical depth (unitless)


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Optical Depth

Optical depth measures the attenuation of light

tu= 1 at s= MFP

The light we see from an optically thick source

was emitted at tapproximately 1


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Radiative Transfer

T1

t1>> 1

T2

t2

What does the observer see?

-assume that the background

cloud is opaque (t1>> 1)

-assume both clouds are uniform

)1)(()()( 21

t

t

t

= eTBeTBF

This equation has two simple limits

Optically Thin (tu


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Optically Thick (tu>> 1)

many scatterings through the cloud

one or fewer scatterings through the cloud on average

)(]1)[()( 21 TBTBF ttt =

contribution from

background cloudcontribution from

foreground cloud

)()( 2TBF t =

since the foreground cloud is optically thick,

all the contribution is from that cloud


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Stel lar Spectra:

The spectrum of a star forms in its

atmosphere

The temperature in the atmosphere is

stratified

The emission at any temperature is a

blackbody (for an optically thick source)

The opacity is a function of wavelength


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At each wavelength, t=1 corresponds to a

different depth in the atmosphere and thus adifferent temperature

The opacity in a line is much higher than in a

continuum In a line, we see to a very shallow depth in the

atmosphere

The Solar Spectrum (from Kitt Peaks McMath-Pierce Solar Telescope):

296013000 angstroms


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g

Based upon the previous image, does


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p p g

temperature in the Sun increase or

decrease with height in the atmosphere?

Temperature decreases with heightbecause the lines

(which are formed higher up) are darker than the continuum

and thus are emitted from a cooler region.

This allows us to probe the temperature of the sun as a

function of depth.

Solar Spectrum Trace:


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notice the different linewidths in different lines

and the strong Calcium H & K lines

Ca K Ca H


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Solar L imb Darkening

Sun is brighter in center

than at edges. Why?

S l L i b D k i


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At center you see down

to a certain depth at t~1

At edge you only see

down to a shallower

depth (lower

temperature) at t~1

Solar L imb Darkening


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Terminology

Surface brightness is synonymous with temperature

The continuum

has a TBof

5,000K

The line has a TB

of 10,000K


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Terminology: Radio Astronomy

Radio astronomers often plot spectra as TB

vs

TBis a physical measurementbut only for thermalprocesses (that are

optically thick)

Why is this terminology most appropriatefor radio astronomy?

radio astronomy is well into the R-J tail


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Terminology: Radio Astronomy

0.1 m

The CMB

S t


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Spectrum

Summary & Goals: Oct 4


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Discuss observing proposals

What we covered:Review EM basics, atomic structure basics

I ntro to gratings & spectrographs: Grating equation

Black-bodies. F luxes, exposure time estimation

The H atom & its spectrum

Einstein A & B coeff icients; radiative & coll ision rates

Radiative transfer & spectral l ine formation

To be covered by F ield Trip:

Saha equation: ionization of atoms into successive stages of

ionizationStel lar classif ication basics

Nebula ionization & excitation: the roles of UV

Spectrograph design: optics, matching R, pixels, and seeing

Radio astronomy

Ionization Balance: (Saha formula)


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Each element has an ionization potential for eachevery electron: Roman numeral is number

of electrons lost + 1

Netural H = HI

Ionized H = HIIMolecular H = H2

HI13.6 eV

HeI24.58 eV, HeII54.416 eVCI - 11.26 eV, CII25.14 eV, CIII47.89 eV, CIV64.49 eV,

CV392 eV, CVI490 eV

OI13.6 eV, OII35.11 eV, OIII54.9 eV

Ca XXI5,469 eV

Ionization stage


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Temperature

I II III IV V

R

elativeabun

dance

Saha Formula


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Saha Formula

Electron density

Next ionization stage density

Previous ionization stage density

Partition function (# of states)Ionization potential

Stellar Spectra: Temperature,


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Ionization state

Dominant features in spectra of stars


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Ionization Balance: Ionized nebulaeHII regions planetary nebulae: UV


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HII regions, planetary nebulae: UV

Supernova remnants: shocksIonization produced by:

- UV to X-ray radiation fields:

stars, white dwarves, neutron stars, accreting WDs, NS,

and black holes

- Collisions: Shock heated gas

Recombinations: Electrons re-combine with ions

Stromgren (photo-ionization equilibrium): HII regions

Q = (4 p/3 r3ne2aBQ = Lyman continuumluminosity (~1049photons/sec for O7 star)aB = 2.6 x 10-13 cm3/sec (Recombination coeff. for H at 10,000 K)

Thors helmet:

NGC 2359


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NGC 2359

HD 56925


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S106 star forming region in Cygnus

(Subaru telescope)

Proto-planetary & Planetary Nebulae


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Orion A:

- Outflows up to

30 l !


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30 pc long !

HH 1/2

M42

HH34HH 131

YSOs near massive stars: UV photo-ablation of disks

irradiated jets


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d253-535 in M43


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HH 46/47

HH 46/47


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HST 1997 - 1994

HH 46/47


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HST 1997 - 1994

Stromgren radius of an HII region:


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Lyman continuum luminosity of O, B stars:

O3: L(LyC) = 1.0 x 1050 photons s-1 T = 60,000 KO5: L(LyC) = 4.7 x 1049 photons s-1 T = 48,000 K

O7: L(LyC) = 6.7 x 1048 photons s-1 T = 35,000 K

O9: L(LyC) = 1.7 x 1048 photons s-1 T = 32,000 K

B0: L(LyC) = 4.7 x 1047 photons s-1 T = 30,000 K

B3: L(LyC) = 4.7 x 1045 photons s-1 T = 20,000 K

n = 1000, O5 star:

L(Lyc) ~ n2r3aB => r ~ [L(LyC) / n2aB]

1/3

5.6 x 1018(cm)

1.8 pc

Photo-ionization equilibrium

(in-class exercise)


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(in-class exercise)

Consider an O7 star that emits 1049Lyman continuum

photons per second which is embedded in a uniform

density cloud with n(H) = 1 cm-3.

- What is the Stromgren radius?

- What is the massthat is ionized?

- How would these answers change ifn(H) = 104cm-3

HII (ionized nebulae) cooled and traced

by trace elements & ions


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Many forbidden transitions have DE ~ 2 eV (visible)long life-times, low decay rates (Einstein A coefficients)

Collision rate: Rcoll= nn ~ 102cm-3

s ~ 1015cm2 (for atoms. Depends on v for ions)v ~ (kT / mm)1/2 (sound speed ~ 10 km/s for H

at 10,000 K)R ~ 10-7sec-1 (1 collision every 107sec)

Collision rate ~ decay rate => each ion can radiate

Thousands of times before recombining => bright line

Some common transitions in ionized nebulae:


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[SII] 6717/6731 A (density tracer)

[NII] 6748/6784 A

Ha 6563 A[OI] 6300/6363 A

[OIII] 5007 A

[OII] 3729/3726 A

Long-slit:


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Spectrum

of a planetary

nebula

Slitless: No entrance aperture


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Objective prism (slitless) spectra:

Planetary nebula M57 (Ring nebula)



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Why `forbidden emission lines are bright


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13.6 eV

~2 eV

9.2 eV

Photo-ionization =>

recombination

Ha

Collisional

Excitation

~3/2 kT = 1.3 eV @104K

-3

1.4

Measuring nebular density using [SII] lines


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[SII]I(6717/6731)

1.0

0.6

101 102 103 104

Density (cm-3)

O star embedded in semi-infinite wall near edge:


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n(H)

d

Q = L(LC) = 1050gs-1

O star next to an infinite wall of hydrogen:


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d

Q = L(LC) = 1050gs-1

Three problems:Star with a wind spherical cloud star in a pipe


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p p p

Ionization Balance: (Saha formula)


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Each element has an ionization potential for each

every electron: Roman numeral is number

of electrons lost + 1

Netural H = HI

Ionized H = HIIMolecular H = H2

HI13.6 eV

HeI24.58 eV, HeII54.416 eVCI - 11.26 eV, CII25.14 eV, CIII47.89 eV, CIV64.49 eV,

CV392 eV, CVI490 eV

OI13.6 eV, OII35.11 eV, OIII54.9 eV

Ca XXI5,469 eV

I II III IV V

Ionization stage


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Temperature

R

elativeabun

dance

Saha Formula


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Electron density

Next ionization stage density

Previous ionization stage density

Partition function (# of states)Ionization potential


Ionization state


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Ionization state

Dominant features in spectra of stars


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Wolf-Rayet stars:

> 60 Solar mass, post-main sequence


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WR 124


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2006 APO Field Trip


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What to Bring:

- Pack light (like carry-on on an airplane)- Jacket, hat, gloves (prepare for cold near freezing)

- Flashlight

- Cash for food (supermarket + stops during drive)

- Personal items

Where:

- Meet at Circle at NW corner of Benson @ 9:00 AM Monday

30 Oct (be early!)- Need two volunteers with sleeping bags for Mon night

(Socorro)

- Return Friday (3 Nov) in the evening.

2006 APO Field Trip


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Itinerary:

- Monday: Drive from Boulder to Socorro, NM (9 - 10 hrs)- Tuesday: Meet Debra Shepherd at NRAO ~ 8:30 AM

Drive to VLA site (1 hr)

Tour VLA

Return to Socorro - have lunch

Drive to APO (4 hrs) & shop for food

Settle in to dorm rooms / houses

Observe till 1:00 AM (If we are late, remote

observers will operate remotely from Boulder)

- Wed: PM tour of NSO (?) + cook dinnerObserve all night

- Thurs: Sleep during day / observe first half

- Friday: Rise at 8:00 AM, drive back (10 - 11 hrs)

Project / Observing Summary

i


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Itinerary:

Tues (first half) M17 LBV, Ceph A DIS new high redHyades WDs (Audrey, Ward, Nate)

Wed (whole night) Comet Swan (Corey, Julia, Tedd)

Eyepiece on Moon etc.

Metallicity

QSO outflow (Max)HL/XZ Tau (Alexi, Courtney, Carlee, Beau)

DIS new high red / eyepiece / DIS / SpiCam / eyepiece (dawn):

Orion, NGC1068, Saturnthurs (first half) APOLLO laser

finish projects as needed.


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Atacama Large Millimeter Array:

Sajnantor Chile, ~ 64 12 meter dishes

Baselines: 150 meter to 10 km


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Baselines: 150 meter to 10 km

ALMA site: Sajnantor Chile,

Elevation ~ 5,000 meters!


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observations & instrumentation ii: spectroscopy

Documents