I. THE EARTH’s ATMOSPHERE
for groundbased observations the Earth’s atmosphere is
important in many ways;
at optical/IR wavelengths:
 1 absorption + scattering extinction
 2 refraction + atmospheric dispersion
 3 background radiation
 4 turbulence  seeing and scintillation
II.1 Extinction
For plane-parallel layers, no refraction, one l, in zenith:
dI / I = -k(h) dh k = extinction m-1
 I / I0 = e-K with K = extinction for atmosphere at zenith
at zenith distance z: I / I0 = e-Ksecz
in magnitudes: m – m0 = -2.5 log (I / I0) = 1.086 K secz
in reality: curvature atmosphere + refraction
K depends strongly on l !
 use airmass M(z) instead of secz for z > 40°
Table shows I / I0 = e-K.M(z) for two
wavelengths in visual/near UV:
cosine
rule in
D ZPS:
sec z = (sin d sin f + cos d cos f cos H)-1
Detection Techniques ’09 – JWP – p7
Main extinction contributions in the classical optical region (l ~ 300-1000 nm):
1) Rayleigh scattering by air molecules;  l –4
2)
Absorption by O3; strong l-dependence especially in near-UV
3)
Scattering by aerosol particles; not very l-dependent
Local variations in aerosol contributions are very large, even
between different observatories at same elevation:
The very steep cutoff at ~ 300 nm is partly
due to l –4 of Rayleigh scattering, but the O3
absorption dominates:
limit for
groundbased
photometry
Sun no
longer
detectable
from H. Tüg,
thesis, 1978
Strong O3 absorption is remarkable in view of
small O3 amount: 3 mm at standard (T,P);
equivalent of whole atmosphere is 8.5 km.
For l < 230 nm O3 opacity drops, but
atmosphere remains opaque because of
Rayleigh scatt. + absorptions by O2 + N2.
Detection Techniques ’09 – JWP – p8
infrared extinction is dominated by molecular
absorptions, mostly of H2O and CO2
 observations are only possible in a number of
‘windows’, designated by I, J, K, L, M, N, Q
H2O absorption
depends strongly
on humidity 
highly variable !
requires high
and very dry
observatory sites
mm
Detection Techniques ’09 – JWP – p9
II.2 Refraction and atmospheric dispersion
Zenith distance as observed from the ground is always < z outside atmosphere.
For precise calculation of refraction one needs:
 curved atmosphere geometry
 refractive index along line of sight  complex: depends on
P,T and composition (humidity !) as function of height
In the visual a good approximation till z = 75° is:
(z – z)550 = (60.4“).tan z - (0.064“). tan3 z
(l = 550 nm, P= 1013 bar, 0° C)
Refractive index n is function of l  dispersion  image becomes spectrum
(z – z) l differs by factor (nl –1)/(n0 –1) with respect to (z – z)550
Typical values in visual l –range:
nair at 550 nm:
1.000293
l (nm)
300
400
500
600
800
(P= 1013 bar, 0° C)
(nl –1)/(n0 –1)
1.047
1.014
1.001
0.995
0.989
In the IR strong H2O absorption
causes deviations from the standard
Edlén formula (Mathar, 2004)
Red: standard
(n-1) formula
Blue: including
effect of H2O
Detection Techniques ’09 – JWP – p10
II.3 Atmospheric background radiation
Main contributions to atmospheric sky background
at dark site:
• twilight at sunset/sunrise
• moonlight
• airglow and aurora
• thermal emission
It is sometimes convenient to
express astronomical surface
brightness in ‘S10 units’:
S10 = equivalent number of 10th
magnitude stars per (deg)2
twilight
sky brightness in zenith at l = 550 nm
(in S10 and mag/(arcsec)2) between
sunset and ‘end of astronomical
twilight’ (hSUN = -18 °) :
without Moon,
Zodiacal light
Ex.: resolution human eye: ~1arcmin  at hSUN =0° sky brightness
per res. element is –1.7mag, i.e. brighter than brightest stars
moonlight
average brightness of
moonlight as a function
of lunar phase,
in mag/(arcsec)2
for V and B
Q. : can you explain
this difference ?
•
Detection Techniques ’09 – JWP – p11
Q. : why this change ?
•
airglow and aurora
airglow (nightglow) is line-emission that arises in the upper atmosphere
from photo-ionization/dissociation and photo-chemical reactions driven by
UV sunlight
nightglow is strongly variable,
most reactions involve O/O2/O3, N/N2, Na, H2O, OH:
both spatially and in time
Ex.: all-sky 6300 Å maps made at
Haleakala (Hawaii), 15min intervals
figures from Roach & Gordon:
‘The Light of the Night Sky’ (1973)
Detection Techniques ’09 – JWP – p12
many different reactions contribute to nightglow:
tables from
Roach &
Gordon:
‘The Light of
the Night Sky’
(1973)
Detection Techniques ’09 – JWP – p13
nightglow lines can be very bright (OI 5577 Å, OH bands in near-IR )
 stay away from them !
near-IR nightglow is dominated by
OH emission (rotation-vibration bands)
a forest of emission lines, but between these lines
the sky is very dark and the lines are narrow
 higher spectral resolution (R > 104) helps !
atmospheric radiation quantities are
frequently expressed in Rayleigh units:
at given l 1 R  4p . (surface brightness B)
with B in units of (106 quanta) . cm-2. s-1.sterad-1
relation to S10: 1 R.Å-1 = 227 S10 (vis)
example of OH lines in H-band:
Table + fig. From
Maihara et al. PASP
105, 940,1993
Detection Techniques ’09 – JWP – p14
Principal upper atmosphere emissions
(Roach & Gordon 1973)
aurora is transient emission driven by high-energy particles from the Sun
mostly a polar phenomenon
 usually unimportant for observatories at latitudes < 40°
when active, aurora can be much brighter than nightglow
most prominent auroral lines:
OI 5577 Å, 6300 Å, Ha 6563 Å
Detection Techniques ’09 – JWP – p15
for l > 2.5 mm thermal radiation
from the atmosphere becomes
the dominant sky background
(hence: the ‘thermal infrared’)
Teff(atmosphere)  250 K
 peak at l  12 mm,
but note: atmosphere is
not a black body !
RADIANCE W cm-2 sr-1 m-1
thermal atmospheric emission
Sea level
Kirchhoff: in thermal equilibrium
emissivity = absorptivity
 IR sky emissivity is ‘mirror image’
of
IR atmospheric transmission curve
this
hits us twice: katm. high

large fraction of source photons removed
and
high sky background
consequence for ground-based observations in the
thermal IR: nearly all astronomical sources are
very faint w.r.t. sky background
Detection Techniques ’09 – JWP – p16
Ex.: a Lyr at l = 20 m:
Fn = 12 Jy  Fl = 5.8 x10-14 W.m-2.m-1
this was outside atmosphere  on the ground: 1.6 x10-14 W.m-2.m-1
sky radiance:  10-5 W.cm-2.sr-1.m-1 = 2.4 x10-12 W.m-2.m-1.(arcsec) –2
so if PSF = 1arcsec (diffraction-limited 8-m telescope at 20 m):
Fsky = 150 x F(a Lyr)
120 mJy source (m20m = 5mag): Fsky = 15000 x Fsource !
solution:
differential
measurements
by means of
‘chopping +
nodding’
Example:
VLT + VISIR
Q-band
spectroscopy
medium
resolution:
R 1500
18.2
FRAME TYPE 1:
SKY+STAR
no chop-subtraction
star invisible w.r.t. high sky level
FRAME TYPE 2: STAR ONLY
removed by subtracting
chopped/nodded frames
nod-pos. 1 
A
B
6” A’
slit length 32.5”
B’
mm
l
18.6
Detection Techniques ’09 – JWP – p17
sky
 nod-pos. 2
NB: all ‘raw’
data, no
flat-fielding
more noisy
horizontal
bands:
drop in S/N at
strong sky line
clusters
-
+
-
3 chopped/nodded spectra of star HD4128 (F18m = 25 Jy)
strong absoptions due to sky lines, but good S/N (> 100)
in clean windows
comparison with sky background outside atmosphere
outside atmosphere the optical/IR sky brightness is determined by:
• dust particles in inner solar system  zodiacal light
• integrated light of faint stars and galaxies
zodiacal light is strongly concentrated
towards Sun and ecliptic plane:
in the range l = 3-70 m
thermal emission from
zodiacal dust dominates
the sky brightness
Gegenschein at 180°
zodiacal light in the visual: scattering
note surface brightness values in S10
Note scale: 1 MJy/sterad
= 2.35x10-5 Jy/arcsec2
ISO data, Leinert et al. A&A 393,1073, 2002
l (m)
compare
with backgr.
of stars +
galaxies:
Detection Techniques ’09 – JWP – p18
II.4 Atmospheric turbulence
turbulence is caused by temperature fluctuations in the
convectively unstable troposphere (h < 10 km)
thermal conductivity of air is low  DT can live long
DP is smoothed out very quickly (sound velocity !)

a) turbulence cells have (DT, Dr) w.r.t. environment, but DP = 0
b) once formed, these cells can live rather long (> 1 s)
the lightpath is influenced by Dr because (n-1)  r
 turbulence cells work as weak positive/negative
lenses that float with wind velocity through the line of
sight causing 2 effects:
fluctuations in direction  ‘seeing’
fluctuations in brightness  ‘scintillation’
apply ideal gas law (index 0 for ambient quantities outside cell):
(n-1)/(n0 –1) = r / r0 = (PT0) / (P0T), P = P0  Dn = -(n0-1).(T0/T2).DT
at sea level:
T  T0  300K, n0 = 1.000293 (l = 550 nm)  Dn  10-6 .DT
at altitude h:
(nh-1) = e-h/H .(n0-1) with H = scale height atmosphere  8.5 km
 Dn = -10-6 .e-h/H .DT
DT  0.1-1 K, h < 10 km  cells have Dn  10-6–10-7
 direction changes in range 0.1-1 arcsec
Detection Techniques ’09 – JWP – p19
very simple model for optical effects of turbulence cell
assume:
atmosphere is isothermal,
in hydrostatic equilibrium
flat wavefront
scintillation:
DI  h.e-h/H
take turbulence cell at height h,
with diameter L and n = n0+Dn
in figure: Dn > 0 (DT < 0)
 wavefront retarded inside cell
seeing:
Df  e-h/H
L” = L –2h.Df
a) direction changes: seeing
L’ = L[n0/(n0+Dn)]  Df = 2(L-L’)/L =
= 2Dn / (n0+Dn)  2Dn = -2x10-6. DT .e-h/H
this function peaks at h = 0
observations confirm this: turbulence in lower layers has
highest weight for seeing (ground layer and dome seeing !)
for small h and typical DT  1°: Df = 2x10-6 rad = 0.4”
Dtel.  L :
at any time only 1 cell in beam  PSFtelescope
unaffected, but whole image shifts
Dtel > L :
image is smeared into seeing disc
without seeing a diffraction-limited telescope with aperture D >L
has PSF-diam. d = 1.2 l/D (D= 10 m, l = 0.5 m  d = 0.01”)
with seeing: image becomes blob of rapidly moving ‘speckles’ with
overall diam. Df regardless of telescope diam. (diam. speckles: 1.2 l/D)
Detection Techniques ’09 – JWP – p20
b) intensity changes: scintillation
DI/I on the ground is proportional to h. Df, i.e.  h.e-h/H
this function has maximum at h = H = 8.5 km
 scintillation comes mostly from high layers
 seeing and scintillation are usually uncorrelated
also these results are confirmed by observations:
• a typical cell with 10 cm at h=8.5 km has angular diam. 2”
 planets don’t scintillate
( Mars: ~5-15”,Jupiter: ~40-50”, Venus: 15-60”)
• the seeing and scintillation fluctiations are usually uncorrelated
scintillation amplitude + frequency
for small telescope apertures:
scintillation amplitude DI/I
- drops below 100 Hz
- decreases strongly with Dtel
this fits with typical cell sizes (5-25 cm) and
wind speeds at 8-10 km (~100 km/h)
DI/I is largest for D   turbulence cell
(e.g. human eye !)
Dtel > 1m  scintillation becomes negligible
(averages out over many cells)
Detection Techniques ’09 – JWP – p21
more realistic description of wavefront
distortions by atmospheric turbulence
Lit.:
Beckers: Ann. Rev. A&A
Vol.31, 13, 1993
a proper description should take into account:
• atmospheric turbulence in reality is a 3-D field of cells
with a distribution of cell sizes and temperature variations
• the atmosphere is in hydrostatic equilibrium, but not isothermal i.e. T = T(h)
NB: from now on we concentrate on seeing,
as scintillation for large telescopes is unimportant
Kolmogorov (1941): in a turbulent gas flow the kinetic energy
of the turbulence eddies with spatial frequencies f is  f-5/3
this power law holds between scale Lu (‘outer scale of turbulence’
= scale at which turbulence is generated) and Ll (lower scale,
where turbulence dissipates by viscosity  very small)
NB: ongoing
dispute about Lu !
for Paranal:
median Lu  22 m
on the basis of Kolmogorov turbulence Tatarski (1961) developed
the theory that is commonly used in astronomical seeing models
most important parameters:
DT(Dr)  <|T(r +Dr) –T(r)|2> (in K2)
for Kolmogorov turbulence:
DT  Dr  Dn
 structure function for n:
Detection Techniques ’09 – JWP – p22
= variance in T for two points Dr apart
DT(|Dr|) = CT2 .|Dr|2/3
(CT2 = ‘structure constant of T-variations’)
Dn(|Dr|) = Cn2 .|Dr|2/3 where
Cn = 7.8x10-5 .(P/T2). CT (at l = 0.5 m, P in mbar)
fluctuations in n cause fluctuations in phase and amplitude
(amplitude  scintillation: can be neglected for large telescopes)
integrated effect of all phase fluctuations along light path through
atmosphere is equivalent with ‘phase screen’ in front of observer that
makes originally flat wavefronts corrugated
for Kolmogorov turbulence the phase ‘structure function’ at the entrance
the telescope is: Df(Dx) = <|f(x +Dx) –f(x)|2> = 6.88 r0-5/3. Dx5/3 (in rad2)
of
with r0 = the coherence length (= ‘Fried parameter’) :
r0(l, z)
= [0.423 .(2p/l)2.sec z .  Cn2(h)dh]-3/5
= 0.185 .l6/5 . cos3/5z . [ Cn2(h)dh]-3/5
convention: unless stated otherwise r0  r0(l = 0.5 m, z = 0°)
other l: scale l6/5
r0 is related to the ‘isoplanatic angle’ 0 = radius of sky area
wavefronts can be considered as coherent (flat)
0 = 0.341 r0/H
(H = average distance of seeing layer)
seeing from thick layers: H = secz. {  h5/3.Cn2dh /  Cn2dh }3/5
where
0 + wind velocity  characteristic timescale:
t0 = 0.341 r0/Vwind
typical
values at
l = 0.5 m :
r0 10 cm (r0  size of average seeing cell)
 typical seeing-dominated PSF (for r0 < Dtel.): d (FWHM)  l/r0  1”
if seeing layer at h=10 km  0  0.7”
with Vwind = 10 m/s: t0 = 0.003 s, or f0 = 300 Hz
Detection Techniques ’09 – JWP – p23
turbulence often
occurs in 3 domains:
incoming flat wavefront
boundary troposphere-stratosphere
h1=10-12 km wind-shear turbulence (jet streams !)
this limits best seeing at best sites
(when turb. at h2 and h3 negligible): r0
 20 cm  l/r0  0.5”, 0(h1)  1.4”
r0 (h1)
high Vwind (~100 km/h)  t0 2 ms
0(h1)
corrugated wavefront
h
h2  1km
r0 (h2)
0 (h2)
h3 < 30 m
r0 (h3), 0 (h3)
NB: if there is turbulence in the lowest layers this
usually dominates seeing because (n-1)h  rh  e-h/H
‘planetary boundary layer’
turbulence from convection driven
by daily solar heating
if r0 (h2) = r0 (h1) 0 (h2) = 10x 0 (h1)
lower wind speeds  t0  5 ms
surface layer
turbulence
from wind-surface interaction (+
man-made !)
small h  0 (h3) up to few arcmin
t0  10 ms
Detection Techniques ’09 – JWP – p24
II: SYSTEM CALIBRATION:
SETTING UP A SYSTEM OF
SPECTROPHOTOMETRIC
STANDARDS
example: the standards for the Walraven
(V,B,L,U,W) photometric system
Lit.: Pel, J.W., Lub, J., in ‘The Future of Photometric, Spectrophotometric and
Polarimetric Standardization’ ASP Conf. Ser. #364, p.63, Ed. C.Sterken, 2007
the Walraven 5-channel photometer was one of
the few simultaneous multi-band photometers
that operated in the period ~1960-’90
other example: Strömgren (uvby) photometers
big advantage of simultaneity:
extinction variations are ~ ‘gray’
 high precision in colours
NOTE: 0.01m colour errors 
significant and correlated errors in
(IS reddening,Teff, logg, [Fe/H], age)
calibration ‘from scratch’
was done twice:
1958 by Th.Walraven during start
of observations at the Leiden
Southern Station in S.-Africa
1979 by J.Lub + J.W.Pel after
move of telescope +instrument
to ESO,Chile because of:
• different atmosphere+elevation
• new photomultipliers
• new telescope mirror coatings
Detection Techniques ’09 – JWP – p48
strategy:
1) select ~ 20 suitable stars in ring
around the sky close to Dec= fgeogr
2) observe star pairs when
D(secz) ~ 0 within Dt < 10 min
 1st approx. of extinction coeff.
are OK for accurate D(mag)i in all
channels, detector gain drifts drop
out due to small Dt
3) repeat pair observations many times
and take averages  accuracy of
D(mag)i per pair: ~0.001mag
4) make least-squares solution for
whole network of pair differences
residuals now <0.0005mag !
5) check residuals for systematics
(dependence on brightness, colour,
position, time….)
6) all magnitude differences (flux ratios) now accurately known, but zeropoints still arbitrary
 define one primary standard as zeropoint
in this case: HD144470 (O-star so SED  BB )
VW (144470) chosen to match VUBV, all four Walraven colour indices defined as zero
7) standards can now be used for determination of extinction coeff. + instrumental zeropoints
for all nights  connect all secondary standards to the ‘ring standards’
8) connect the whole set of standards to stars with absolute flux calibrations
this gives relations [(VBLUW) magnitudes  physical flux densities]
Detection Techniques ’09 – JWP – p49
9) external verification
the ‘ring solution’ method produced standards with high internal
precision, but there might still be hidden systematic errors
for checks with external data, a set of ~2000 stars with
multiple high-quality VBLUW obs. was compared with
data in the Johnson (U B V), Strömgren (u v b y) and
Geneva (U B1 B2 V1 V2 G) systems
comparison was possible for
two transformable quantities:
VW  VJ  VG  yS and
(V-B)W  (B-V)J  (B2-V1)G  (b-y)S
result:
a) large scatter and systematic RA-dependent differences in comparison with VJ,(B-V)J
 no surprise: UBV photometry is very inhomogeneous
b) systematic RA-trends in comparison with VG, (B2-V1)G  very unexpected !
c) smallest scatter and no systematics in comparison with yS, (b-y)S
 Walraven and Strömgren systems OK !
Detection Techniques ’09 – JWP – p50
additional checks with
data from space
1) comparison between VW and
HIPPARCOS Hp magnitudes:
no systematic trends as a
function of (colour, RA, DEC)
 both photometric systems are
‘flat’ to level of 0.001 mag
2) comparison between WW and
IUE spectrophotometry for SN 1987a (LMC)
 excellent agreement
Note: this now includes the absolute
flux calibration, for l close to the UV
atmospheric cut-off !
Detection Techniques ’09 – JWP – 51