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No. 55 - March 1989
Figure 1: Fa/se-colour images of He2-104 in the light of Ha 6563, [NIIJ 6584 and [SIIJ 6731.
They have been taken with the adapter/GGO at the 2.2-m telescope at La Silla in March 1988.
The exposure times are 30 minutes for the Ha and [NIIJ images and 1 hour for the [SIIJ frame.
Oata reduciion was done at ESO, La Silla, and at the JAG, Hawaii.
The Southern Crab
Heavenly Crabs
The objeet under serutiny in this arti-
eie - a suspeeted symbiotie star with
the prosaie name of He2-104 - looks far
more like a erab than the Grab nebula. A
quiek look at Figure 1 will eonvinee you
that the name Southern Grab is most
suitable and fully deserved! Our seareh
for sueh objeets was started by selee-
tion of speeial eases in the adapted
emission line ratio diagram used by
Sehwarz (1988). Four possible prato-
planetary nebulae were found of whieh
He2-104 is by far the most speetaeular
A reeent literature seareh on He2-104
yielded 36 referenees. The objeet was
elassified by 17 authors as a planetary
nebula (PN), by 14 as a symbiotie star,
and 5 simply referred to it as an emis-
sion-line star. The spatial extent was in
all eases reeorded as less than 5". Re-
eent distanee determinations based on
the Gudworth method gave 6 to 7 kpe.
Near IR photometry was used to
elassify the objeet as a possible symbio-
tie star eontaining dust (Allen, 1984) and
a reeent IR photometrie period of 400
days has been found by Whiteloek
(1988), indieating a Mira-Iike star.
No emission-li ne variability has been
observed in the speetrum whieh shows
lines up to [FeVII], [ArV] and Hell but
also low-exeitation lines like [S 11], [N 11)
and [01). Generally, the speetrum is that
of a dusty symbiotie star and is quite
Similar to that of H1-36 (Allen, 1983).
He2-104 is in the IRAS point souree
catalogue as 14085-5112 and has flux-
es on the order of 10 Jy up to 100 mi-
New Observations
Inspeetion of Figure 1 shows that the
objeet extends over about 75" or is 15
times larger than previously thought.
The three images in the light of the Ha
6563, the [N 11) 6584 and the [S 11) 6731
Iines elearly show that the strueture is
different in each of the emission lines.
Exeitation and therefore density and
temperature vary throughout the nebula.
We interpret the object as folIows.
The "Iegs of the Crab" are bubbles
formed by the red giant primary stellar
wind which due to the presence of a
dust torus cannot escape in the equato-
rial direction. The Herbig-Haro type
blobs at either end are moving shocks
due to a highly collimated wind from the
hot secondary which is surrounded by a
thick accretion disko For more details on
this kind of scenario, see Morris (1987).
High resolution spectra taken at the
2.2-m telescope at La Silla show that
the radial velocity of the central abject is
-139 km per second. The blobs have
velocities of -36 and -235 km per sec.
indicating that they are moving away
from the central star with at least
100 km per second. Since the object
looks like being near the plane of the
sky, the space velocities could be as
high as 400 to 600 km per second.
Photometry in the visual and IR com-
bined with the IRAS data have allowed
us to determine the energy distribution
of He2-104. From this curve the
bolometric magnitude has been derived
and when the fluxes are de-reddened
with a measured A
of 1.5 magnitude
and assuming that the central star has a
bolometric magnitude of -2 we arrive at
a distance of 1 kpc. This places the ob-
ject much nearer than previously
thought and this is also borne out by
using arguments on the size and mass
limits of the ejected nebula by compar-
ing with typical values for PNe.
Other independent estimates of the
distance of He2-104 can be obtained by
comparing the absorption with the 21-
cm hydrogen flux in the same direction.
A distance of more than about 1-2 kpc is
also made unlikely by considering that
with a galactic latitude of 10 degrees,
the object would be out of the plane by
about 1.2 kpc, assuming a distance of
7 kpc.
With our derived distance of 1 kpc the
Herbig-Haro objects are at a distance
from the central object of 0.2 pe giving,
with the measured expansion velocity of
100 km per second, a dynamical age of
the nebula of 2,000 years. Since we
have ignored any projection effects on
the radial velocity, this represents an
upper limit to the age which could actu-
ally be as low as 300-500 years.
The Future
Clearly, more work needs to be done
on this fascinating object: we have
planned SEST observations to look
for the presence of an SiO maser, high
resolution spectroscopy to determine
the velocity structure in detail and pho-
tometric imaging to find the temperature
and density structure of the nebula.
IR spectrophotometry could tell us
something more about the central star(s)
since nothing is detected in the visual.
Finally, IR imaging with the new ESO IR
camera, IRAC, could provide informa-
tion about the dust distribution and tem-
perature in the object.
In the near future we hope to spend a
lot more time with our southern Crab!
Allen, D.A.; Mon. Not. R. Astr. Soc. 204, 113
Allen, D.A.: Proc. Astr. Soc. Austr. 5, 369
Morris, M.: Publ. Astr. Soc. Pac. 99, 1115
Schwarz, H.E.; in lAU Coll. 103, 123 (1988).
Whitelock, P.A.: in lAU Coll. 103, 47 (1988).
Key Programmes on La Silla: First Allocations
H. VAN DER LAAN, Director General, ESO
Class Number of proposals Comparison
1. Galaxies and clusters 14 B.B
2. Quasars, Seyferts, radio galaxies B 6.2
3. Magellanic Clouds 3 3.4
4. Interstellar matter 3 5.9
5. Star clusters and galactic structure 4 3.8
6. X-ray sources 1 1.3
7. Stars 7 10.4
8. Solar system 1 1.8
9. Miscellaneous 1 0.4
With the Observing Programmes
Committee's recommendations in hand,
early December last year, all information
was available to allocate and plan the
observations for Period 43 (normal pro-
grammes) and Periods 43 and 44 (Key
Programmes). The scenario of prepara-
tions and decisions has precisely
followed the intentions outlined in my
original article in the Messenger of
March 1988 (No. 51). Here the commu-
nity is briefly informed about the pro-
cedures, the proposals and the alloca-
tions. Investigators of Key Programme
Proposals successful in this first round
will themselves describe their research
plans in brief Messenger articles, start-
ing with three such contributions in this
issue and to be continued. Later in this
article there is an overview of Key Pro-
gramme time committed and time ex-
pected to be available for subsequent
rounds, with the 15th of October 1989
as the next deadline.
Proposal Evaluation
There were 42 proposals submitted
for the 15 October deadline last year.
These were grouped according to the
normal OPC classification (see Table 1).
The column "Comparison" in the table
gives the relative numbers of proposals
per cl ass, normalized to the same total,
received for periods 38 to 42 incl., more
than 1,750 proposals in total. The differ-
ence, where significant, is interesting
and is probably attributable to at least
two factors. First, the Key Programmes
concern primarily the use of the bigger
telescopes (1.5-m and larger); secondly
it has been emphasized, in written and
verbal presentations and in discussions,
that Key Programmes afford an oppor-
tunity to gently re-orient European as-
tronomy towards fields now relatively
underdeveloped in our community but
obviously crucial for the VLT era. This
emphasis appears to provide an ap-
propriate stimulus, resulting in more
proposals of class 1 + 2 and fewer of
class 4 through 7.
For each class of programmes two
referees external to the OPC, mostly but
not exclusively from ESO member
states, and one OPC member were
asked to evaluate and grade the pro-
grammes as weil as to rank them within
their class. A special form designed for
the purpose asked the referees to justify
their grade and rank in prose. Without
exception everyone approached to con-
tribute their expertise and time towards
Class 1:
"A redshift survey of galaxies with z :s 0.6"
Allocation: 60 nights at the 3.6 m plus NTT
Investigators: Oe Lapparent (Inst. d'Astrophysique, Paris)/Mazure (Univ. du Languedoc)/Mathez, Mellier (Obs. du Pic du Midi et de
Class 1:
"Towards a physical classification of early-type galaxies"
28 nights at the 3.6 m and
34 nights at the 2.2 m or 1.5 m Danish
Investigators: Bender (Landessternwarte Heidelberg)/Capaccioli (Oss. Astronom. Padova)/Nieto (Obs. de Toulouse)/Macchetto
(STScl Baltimore)
Class 1:
"A search for dark matter in elliptical galaxies"
Allocation: 9 nights at the 3.6 m
16 nights at the 2.2 m
Bertola (Ist. di Astronomia, Padova)/Bertin (Scuola Normale Superiore, Pisa)/Buson (Oss. Astronom. Padova)/
Danziger (ESO-HQ)/Dejonghe (Koninklijke Sterrenwacht Brussels)/Sadler (Anglo-Australian Obs.)/Saglia (Scuola
Normale Superiore Pisa)/Vietri (Oss. Astrofisico Firenze)/de Zeeuw (Inst. for Advanced Studies Princeton)/Zeilinger
(Ist. di Astronomia, Padova)
Class 1:
"Identification of high redshift galaxies with very large gaseous halos"
25 nights at the 3.6 m
Bergeron (Inst. d'Astrophysique, Paris)/Cristiani (Asiago Obs.)/Pierre, Shaver (ESO-HQ)
Class 1:
"Structure and dynamical state of nearby clusters of galaxies"
Mazure (Lab. d'Astronomie, Montpellier)/Katgert (Sterrewacht Leiden)/Dubath (Obs. de Geneve)/Focardi (Dpto.
Astron. Bologna)/Gerbal (Obs. de Meudon)/Giuicin (Dpto. Astron. Trieste)/Jones (NORDITA, Copenhagen)/Lefevre
(Canada-France-Hawaii Tel. Corp.)/Moles (Obs. Grenada)
Class 1:
"Peculiar motions of rich clusters of galaxies"
Rhee, Katgert (Sterrewacht Leiden)
35 nights at the 3.6 m and
5 nights at the 1.5 m Danish
for the combination of the above mentioned two Key Programmes.
Class 2:
"A study of the most distant radio galaxies"
43 nights at the 3.6 m or the NTT and
34 nights at the 2.2 m
Miley (Sterrewacht Leiden)/Chambers (STScI, Baltimore), Hunstead (Univ. Sydney)/N. N. (La Silla Postdoc, beginning
period 45)/Roland (Inst. d'Astrophysique, Paris)/Röttgering (Univ. Leiden)/Schilizzi (Radiosterrenwacht, Dwingeloo)/
Macchetto (STScl Baltimore)/N.N. (Leiden Postdoc beginning Sept. 1, 1989)
Class 2:
"Gravitational lensing: Quasars and radio galaxies"
54 nights at the 3.6 m or the NTT and
48 nights at the 2.2 m and
9 nights at the 1.5 m Danish
Surdej (Inst. d' Astrophysique, Cointe-Ougree)/ Arnaud (Canada-France Hawaii Tel. Corp.)/Borgeest (Hamburger
Sternwarte)/Djorgovski (CALTECH)/Fleischmann (Phys. Inst. der Univ. Erlangen-Nürnberg, Erlangen)/Hammer (Obs.
de Meudon)/Hutsemekers (ESO-La Silla)/Kayser (CITA, Univ. of Toronto)/Lefevre (Canada-France Hawaii Tel. Corp.)/
Nottale (Obs. de Meudon)/Magain (Inst. d'Astrophysique, Cointe-Ougree)/Meylan (ESO-HQ)/Refsdal (Hamburger
Sternwarte)/Remy (ESO-La Silla)/Shaver (ESO-HQ)/Swings (Inst. d'Astrophysique, Cointe-Ougree)Nanderriest (Obs.
de Meudon)/Van Drom (ESO-La Silla)/Veron-Cetty, Veron (Obs. de Haute-Provence)/Weigelt (MPI für Radioas-
tronomie, Bonn)
Class 2:
"A homogeneous bright quasar survey"
40 nights at the 2.2 m
50 nights at the 1.5 m
40 nights at the 1.0 m and
40 nights at the Schmidt telescope
Barbieri (Oss. Astronomico, Padova)/ Andreani (Physics Dept. Roma)/Clowes (ROE Edinburgh)/Cremonese (Oss.
Astronomico, Padova)/Cristiani, Gemmo (Asiago Obs., Padova)/Gouiffes (ESO-La Silla)/Iovino (Obs. Merate)/La
Franca (Asiago Obs., Padova)
Class 3:
"Coordinated Investigation of selected regions in the Magellanic Clouds: Population, Structure, Evolution"
36 nights at the 3.6 m
8 nights a the NTT and
37 nights at the 2.2 m
Oe Boer (Sternwarte Univ. Bonn)/Azzopardi (MarseiJle)/Baschek (Landessternwarte Heidelberg)/Dennefeld (Inst.
d'Astrophysique, Paris)/Israel (Sterrewacht Leiden)/Molaro (Trieste)/Seggewiss (Sternwarte Hoher List, Daun)/Spite
(Obs. de Meudon)/Westerlund (Uppsala)
Class 5:
"Radial velocities of southern late type HIPPARCOS stars"
150 nights at the 1.5 m Danish
Mayor, Duquennoy, Burki, Grenon (Obs. de Geneve)/Imbert, Maurice, Prevot (Marseille Obs.)/ Andersen, Nordstrom
(Copenhagen Obs.)/Lindgren (ESO-La Silla)/Turon (Obs. de Meudon)
Class 5:
"Astrophysical fundamental parameters of early-type stars of the HIPPARCOS survey"
160 nights at the 1.5 m Danish
Gerbaldi (Inst. d'Astrophys.)/Gomez, Grenier (Obs. de Meudon)/Faraggiama (Univ. di Trieste)/Turon (Obs. de Meudon)
Period 43 + 44 45 + 46 47 + 48 49 + 50 51 + 52 53 + 54
Telescope Number of nights
K.P. Total 40 + 40 60 + 60 60 + 60 60 + 60 60 + 60 60 + 60
3.6 Committed 15 + 31 39 + 44 14 + 22 4+ 4 0 0 0 0
Available - 21 + 16 46 + 38 56+56 60 +60 60 + 60
K.P. Total 15 + 30 60 + 60 60 + 60 60 + 60 60 + 60 60 + 60
ND Committed 14 + 19 20 + 13 16 + 13 10 + 0 10 + 0 10 + 0
Available - 40 + 47 44 + 47 50 + 60 50 + 60 50 + 60
K.P. Total 20 + 20 40 + 40 40 + 40 40 + 40 40 + 40 40 + 40
2.2 Committed 21 + 43 30 + 40 17 + 28 4+ 4 4 + 4 0 0
AvailabJe - 10 + 0 23 + 12 36 + 36 36 + 36 40 + 40
K.P. Total 25 + 25 50 + 50 50 + 50 50 + 50 50 + 50 50 + 50
1.5 E Committed 16 + 24 21 + 24 21 + 24 21 + 22 21 + 22 0 0
Available - 29 + 26 29 + 26 29 + 28 29 + 28 50 + 50
K.P. Total 15 + 15 30 + 30 30 + 30 30 + 30 30 + 30 30+30
1.5 D Committed 19 + 25 19 + 20 16 + 18 15 + 16 15 + 15 0 0
Available - 11 + 10 14 + 12 15 + 14 15 + 15 30 + 30
CAT/1 m/Schmidt: Available as pressure requires.
the refereeing effort responded positive- TABLE 3. Key Programme Commitments - January 1989
Iy. All reports were received in time for
the OPC meeting. This participation is
gratifying and ESO is pleased to ack-
nowledge the referees' very substantial
contribution to the quality of the pro-
gramme evaluation.
The OPC spent one half day to dis-
cuss the Key Programmes and to com-
plete its recommendations for the allo-
cation of telescope time. As is weil
known, allocation of time on the 3.6-m,
the 2.2-m and the Danish 1.5-m has
long been a process of choosing be-
tween good and very good, with the
cutoff line falling among the good. This
was true for these proposals too. Refer-
ence to the special criteria I sketched in
previous articles in this journal (Nos. 51
and 52), the prospects of some space-
based telescopes and the distribution
among the several programme c1asses
all played their part in the final alloca-
The First Key Programme
Table 2 gives the names and the in-
stitutional affiliation of all investigators
for each successful proposal. The first
name is that of the Principal Inves-
tigator/Coordinator. The class and the
programme title are then given, followed
by the allocation, specified according to
telescope to which ESO is provisionally
committed. Provisionally, because only
for Periods 43 and 44 are the allocations
definitive; the affirmation of the further
allocations depends on the evaluation of
progress reports to be submitted to the
OPC each year in October.
Each one of these programmes will be
introduced by one or more of the pro-
posers in abrief Messenger article in the
course of 1989. The authors have been
asked to clearly sketch their strategy,
goals and their programme's potential
for opening new dimensions in Euro-
pean astronomy. 1hope that these arti-
c/es, together with the information in
Table 2, will further extend existing col-
laborations as weil as serve to inform
new proposers as they prepare to com-
pete in subsequent rounds.
The Scope of the Next Round
In late November this year the OPC
will again assess Key Programme ini-
tiatives. The proposal deadline is the
same as for the normal proposals that
pertain to Period 45, viz. 15 October
1989. In Table 3 an overview is given of
total time to be provided, time al ready
committed, and the difference, that is
the time available. The next round
should fill the time for periods 45 and 46,
/eaving a wedge for later periods.
I have not included the CAT, the
Schmidt and the 1.0-m telescope in the
table. It has become clear that these
instruments must also playa role in Key
Programmes, often in conjunction with
the use of bigger telescopes. Since the
pressure on these smaller telescopes is
relatively lower, they will be made avail-
able for Key Programmes as warranted
by proposal pressure.
The first round was a heavy one and
the allocations have been very substan-
tiai, in order to provide a flying start of
the new service. Both the ESO 1.5-m
and the Danish 1.5-m telescopes are
heavily committed to programmes en-
riching the HIPPARCOS mission and
new commitments in the next round that
go beyond periods 45 and 46 will have
to be light if there are to be opportunities
in 1990 through '93. The 2.2-m MPG
telescope proved very much in demand
and is largely committed already for
next year. For the 3.6-m and the ND
there are 124 nights available for Key
Programmes initiatives in periods 45
and 46. Allocations beyond these
periods have to be modest if the
accumulation is not to pre-empt subse-
quent rounds.
In 1990 a steady state of sorts will
have to be attained for the resources
available each double period. Subse-
quent versions of Table 3 will monitor
this development. They provide some
guidance for users to weigh their ambi-
tions against realistic expectations.
Some Fringe Benefits
It is obvious that the process of writ-
ing the rather extensive, often complex
proposals by multinational teams, was a
big chore and one that could hardly
have been achieved without electronic
mail networks. Even for the 70 % of
proposals that did not receive time in
this round, that effort is not wasted and
may weil serve to guide collaborative
research of a more modest scope in the
near future. In any case, one disappoint-
ment need not deter another try. A con-
versation with the national representa-
tive in the OPC mayaiso improve the
next version as will a very critical look at
the transparency of presentation and
the balance between time requested
and results hoped for.
Several successful Key Programmes
have led to small workshops among the
proposers, as they get their act together
for the heavy observing schedules and
reduction labour ahead. In one case I
organized a two day get-together at
ESO Headquarters to induce the teams
of half a dozen responses to the prelimi-
nary enquiry one year ago, to work
together. The major initiative for
Magellanic Cloud research, coordinated
by K. S. de Boer in Bonn and combining
these teams' intentions, was the result. I
hope this coordination will have a semi-
na/ effect upon investigations of the
Clouds that will spill over even beyond
the scope of this far ranging specific
programme. Another example of a fringe
benefit is the renewed interest in the
ESO Schmidt telescope as a research
tool, discussed in a Garching workshop
in early February. The outcome will be
reported in one of the next issues of the
Messenger. One year after the Key Pro-
gramme initiative it may be concluded
that at least one goal: more European
and more ambitious interaction, is being
A Redshift Survey of Galaxies with z ::s 0.6 Using Multi-Object
11. DELAPPARENT, Institut d'Astrophysique de Paris, France
A. MAZURE, Laboratoire d'Astronomie de Montpellier, France
G. MATHEZand Y. MELLlER, Observatoire de Tou/ouse, France
Through the steady acquisition of red-
shifts of galaxies, our understanding of
the large-scale galaxy distribution has
evolved drastically during the past de-
cade. Most recently, the Centre for As-
trophysics (CfA) redshift survey (B =5
15.5) has suggested a new picture of the
galaxy distribution: galaxies appear to
be distributed on thin shells surrounding
vast regions with diameters between 20
and 50 h-
Mpc (Ho = 100 h km S-1
) devoid of bright galaxies (de
Lapparent, Geiler, and Huchra 1986).
This new interpretation of the galaxy
distribution is consistent with the detec-
tion of a 60 h-
Mpc void in Boötes
(Kirshner et al. 1981) and with the alter-
nation of peaks and - 100 h-
wide valleys in deep pencil-beam
probes (Koo, Kron, and Szalay 1987).
Shell-like structures are also detected in
redshift surveys of HI galaxies (Haynes
and Giovanelli 1986).
the CfA redshift survey is
unlque In ItS combination of Jarge angu-
lar and depth, it does not rep-
a fair sampie of the universe: the
slze of the Jargest coherent structures
(- 50 h-
Mpc) is comparable to the
spatial extent of the survey (- 100 h-1
Mpc). Deeper redshift surveys are there-
fore necessary for measuring the mean
statistical properties of the large-scale
clustering. Deep surveys can constrain
the size and frequency of the shells, and
thus complement nearby wide-angJe
surveys which provide information
about the details of the galaxy distribu-
tion within the shells. In an attempt to
satisfy the need for deep surveys, we
have designed a redshift survey pro-
gramme - 15 times deeper in velocity
than the CfA redshift survey.
Our programme aims at the acquisi-
tion of a complete galaxy catalogue to
R - 20.5, with BVR photometry and
low-resolution spectroscopy (ßA. -
10 A). We will start the observations on
the 3.6-m using the ESO Faint Object
Spectrograph and Camera (EFOSC;
Buzzoni et al. 1984) for both imaging
and spectroscopy. We will take full ad-
vantage of the multi-object spectros-
copy capability of EFOSC provided by
the Punching Machine (PUMA), which
allows real time drilling of multiple slits
into aperture plates (Fort et al. 1986). At
R :s 20.5, there are on the average - 8
galaxies per - 3 x 5' field of EFOSC,
and the slit spectra of :s 10 galaxies can
be obtained simultaneously (as demon-
strated during two previous observing
runs). When the ESO Multi-Mode Instru-
ment (EMMI; D'Odorico et al. 1986) be-
comes available on the NTT with the
multi-slit mode (1990), we will switch to
the NTT. The larger field of view planned
for EMMI (7' x 10' with a new-genera-
tion CCD) will result in a time gain of a
factor - 3 for the photometry and a fac-
tor - 2 for the spectroscopy. The total
observing time (60 nights) will be spent
in the proportion of - 1/8 for imaging and
-7fa for spectroscopy.
We will survey a strip with a total area
of 0.4 deg
, in a region near a == 22
, located at high galactic
latitude (b" - -50
). Measurement of
the redshifts of the - 700 galaxies in the
survey region will provide a unique data
base for studying the ciustering on very
large scales. With its effective depth of z
- 0.6 (-1400 h-
Mpc with qo = 0.1),
our survey could intercept along the line
of sight - 45 shells with mean diameter
30 h-
Mpc. The size of the projected
area surveyed by the strip at 1400 h-
Mpc, -6 x 45 h-
, can be com-
pared to the characteristic scales for the
structures in the CfA redshift survey
(mean separation of galaxies in the
shells - a few h-
Mpc -, and diameter
of the shells). This comparison suggests
that the deep survey will sampie the
structures with a sufficient spatiaJ cov-
erage for delineation of the individual
The deep redshift survey will indicate
whether the nearby voids are typical
in size, and whether larger in-
homogeneities exist. The contribution of
large voids dominates the uncertainty in
the mean number density of galaxies.
This uncertainty in turn introduces a
large scatter in the determination of the
two-point correlation function (de Lap-
parent, Geiler, and Huchra 1988), an
important tool for constraining the pow-
er spectrum of primordial fJuctuations at
large-scale (Peebles 1980). In addition,
the survey will provide information on
the spectrum of shell diameter, and
therefore, will tightly constrain the N-
body models for the formation of large-
scale structure (Melott 1987; White et al.
1987; see also Ostriker and Vishniac
1986). The large number of individual
structures sampled by the survey will
also allow to put direct limits on evolu-
tion in the large-scale clustering with
The other major goal of the pro-
gramme is to examine the spec-
trophotometric evolution of field galax-
ies. Catalogues of distant clusters of
galaxies (Butcher and Oemler 1978;
Butcher and Oemler 1984; Dressler,
Gunnn, and Schneider 1985; Lavery and
Henry 1986) suggest a significant evolu-
tion at z - 0.5 in the "blue" population
of clusters. Deep galaxy counts suggest
a similar effect in the field at B 24 (Koo
1986; Tyson 1988), suggesting UV ex-
cess from massive star formation in
galaxies at z - 1-3. The availability of
redshifts for the galaxies in our survey
will allow to recover their intrinsic col-
ours and thus distinguish evolution from
cosmology. We will then be able to esti-
mate the fraction of intrinsicaJly "bJue"
galaxies [(B-V)inlr :s 0.7) per redshift in-
terval. Examination of the spectral
features will yield information on the na-
ture of these "blue" galaxies (see Dress-
ler and Gunn 1983), and thus will pro-
vide a test for the interpretation of the
"blueing" trend as originating from ex-
tensive star formation. A study of the
variations in the fraction of "blue" galax-
ies with the local galaxy density might
also provide some insight into the influ-
ence of the environment on the process
for galaxy formation.
Butcher, H.B., and Oemler, A.Jr.: 1978, Ap.
Butcher, H.B., and Oemler, AJr.: 1984, Ap.
Buzzoni, B., Delabre, B., Dekker, H.,
D'Odorico, S., Enard, D., Focardi, P., Gus-
tafsson, B., and Reiss, R.: 1984, The
Messenger, 38, 9.
D'Odorico, S. et al.: 1986, Proc. SPIE, 627,
Dressler, A, and Gunn, J. E.: 1983, Ap. J.,
Dressler, A, Gunn, J. E., and Schneider, D. P.:
1985, Ap. J., 294, 70.
Fort, 8., Mellier, Y., Picat, J.P., Rio, Y., and
LeliEMe, G.: 1986, Proc. SPIE, 627, 321.
Haynes, M.P., and Giovanelli, R.: 1986, Ap. J.
(Letters), 306, L55.
Kirshner, R. P., Oemler, A. Jr., Schechter,
P.L., and Shectman, S.A.: 1981, Ap. J.
(Letters), 248, L57.
Koo, D.C. 1986, Ap. J., 311, 651.
Koo, D.C., Kron, R.G., and Szalay, A.S.:
1987, in Thirteenth Texas Symposium on
Relativistic Astrophysics, ed. M. P. Ulmer
(Singapore: World Scientific), p. 284.
de Lapparent, V., Geiler, M.J., and Huchra,
J.P.: 1986, Ap. J. (Letters), 302, L1.
de Lapparent, V., Geiler, M.J., and Huchra,
J.P.: 1988, Ap. J., 332, 44.
Lavery, R.J., and Henry, J.P.: 1986, Ap. J.
(Letters), 34, L5.
Melott, A.L.: 1987, Mon. Not. Roy. Astr. Soc.,
Ostriker, J.P., and Vishniac, E.T.: 1986, Ap.
J. (LeUers), 306, L51.
Peebles, P.J. E.: 1980, The Large-Scale
Structure of the Universe (Princeton:
Princeton University Press).
Tyson, JA, 1988, Astr. J., 96, 1.
White, S. D. M., Frenk, C. S., Davis, M., and
Efstathiou, G. 1987, Ap. J., 313, 505.
Towards a Physical Classification of Early-type Galaxies
R. BENDER, Landessternwarte, Heidelberg, F. R. Germany
M. CAPACCIOLl, Osservatorio Astronomico, Padova, Italy
F. MACCHETTO, Space Telescope Science Institute, Baltimore, USA
J.-L. NIETO, Observatoire de Toulouse, France
26 ...,.......,-r-"T"'""-'-""""-"
•• . t I
, .. ... .
20 -.. I ••a ••••
and Möllenhoff 1987, and Jedrzejewski
Bender et al. (1988, 1989) and Nieto
(1988) found that these two "shape-
(a) contain =80 % of all Es, and
(b) isolate objects with distinct physical
More specifically, boxy ellipticals are
frequently radio-loud (the power being a
function of the mass) and permeated by
gaseous "halos" (as inferred from their
X-ray emission), while ellipticals with
pointed isophotes are usually unde-
tected in the radio and X-ray bands
(Fig. 1). There is also a clear segregation
of the kinematical properties between
these two classes: in most cases boxy
ellipticals are flattened by anisotropie
velocity dispersion while "pointed" E's
are rotationally supported (Bender
1988a, Nieto et al. 1988; see Fig. 2).
The nature of "pointed" E's is easily
understood if one assumes (Nieto et al.
1988) that they are internally structured
·1 0 1 2 3
0(4) I ox100
Figure 1: Logarithm of radio lum/nos/ty at
1.4 GHz against isophotal shape parame-
trized by a(4)/a(a(4)/a < 0 indicates boxy
isophotes. a(4)/a > 0 disky isophotes). Below
the dashed line most symbols indicate upper
limits (Bender et al. 1989).
ity dispersion and the mean surface
brightness. The fundamental plane cor-
responds to the cooling diagram plane
of modern theories of galaxy formation,
which are based on the self-regulation
of the dissipative collapse of pro-
togalaxies (Silk 1983, Blumenthai et al.
In contrast to the "global" parameters,
those describing the detailed shape and
kinematics of early-type galaxies (ellip-
ticity gradients, isophotal twists, trends
of radial light profiles, velocity aniso-
tropies, etc.) do not correlate with the
fundamental plane. This lack of correla-
tions has instilled further suspicion that
even the Hubble E class alone is not
physically homogeneous (Capaccioli
1987). lndeed, while all ellipticals share
roughly the same morphological
appearance, they may nonetheless be-
long to genetically distinct families (see
Bender et al. 1988, Nieto 1988, Prugniel
et al., 1989). The consequences of this
hypothesis are more dramatic than hav-
ing the E galaxies populate a "2 + N"
parameter space (e. g. Djorgovski and
Davis 1987, Dressler et al. 1987) since a
physical segregation of E types would
bear directly on our understanding of
the formation/evolution of the different
species of galaxies.
An effective tool for separating E
galaxies into at least two physically dis-
tinct c1asses has been recently iden-
tified with a simple and direct observa-
tional parameter which measures the
lowest-order axisymmetric deviation (a4)
of the galaxy isophotes from perfect
ellipses (Bender 1987). This parameter,
easily extracted from good SIN images
by Fourier analysis, discriminates be-
tween "boxy" (a4 < 0) and "pointed" (84
> 0) isophotes (e. g. Lauer 1985, Bender
Hubble was the first who succeeded
in classifying galaxies within a scheme
of some physical meaning. Although it
soon became clear that Hubble's tuning
fork does not represent an evolutionary
sequence, this essential diagram has
proven to be a powerful tool especially
for the understanding of late-type galax-
ies. On the other hand, the "early-type"
sequence of elliptical (E) and SO galax-
ies is less satisfying, because it does not
seem to reflect a unique sequence of
physical properties. The SO class, al-
though conceived to bridge the gap be-
tween disk- and disk-Iess galaxies, has
often been abused to host ellipticals
exhibiting peculiarities incompatible
with their definition as structureless ob-
jects. For the elliptical galaxies themsel-
ves, "ellipticity" has been found to be
essentially meaningless with regard to
their angular momentum properties, and
shows Iittle, if any, correlation with other
global parameters. This fact became
apparent after the first stellar kinemati-
cal measurements of luminous ellipti-
cals (Bertola and Capaccioli 1975, IIling-
worth 1977); E galaxies are not
necessarily f1attened by rotation and
may have anisotropie velocity disper-
sions (Binney 1978).
A first step towards a physical
classification of early-type galaxies was
the investigation of the correlations be-
tween the c1assical global parameters
(Iuminosity, scale length, projected cen-
tral velocity dispersion, central density,
and metallicity/colour). These parame-
ters were found to form a plane in the
parameter space ("the fundamental
plane"). More specifically, each of these
global parameters can be given as a
function of two independent parame-
ters, namely in general the central veloc-
__ .l.....-__-:
-1.2 -1.0 -0.5 0.0 +0.5 +10 +15
0(4)/0 *100
Figure 2: Anisotropy parameter Nm/a)* against isophotal shape parameter a(4)/a. Boxyellipti-
eals are in most eases supported by anisotropie veloeity dispersions (i. e. Nm/a)* < 1), while
near/y all disk-elliptiea/s are rotationally flattened objeets showing (Vm/a)* = 1. Open eire/es
indieate elliptieals less luminous than MT =-20.5 (Ho =50 km/s/Mpe), while tilled dots indieate
objeets above MT = -20.5 (Bender 1988a).
log (MIM<:»
Figure 3: Speeitie angular momentum J/M (ealeulated trom half-light radius and maximum
rotation veloeity) against mass in so/ar units tor different types ot objeets. Se and Sb ga/axies
are represented by open eireles and squares (adapted trom Fall 1983), disk-E's by erosses,
and boxy E's by filled eire/es (Bender et al. 1989, in preparation).
lished now, the cosmological interpreta-
tion of this finding is still open. The two
classes certainly reflect different histori-
cal pasts in terms of interaction with the
environment, notably at early stages of
galaxy formation.
In particular, it seems that disk-E's,
together with 80's and spirals, belong
to a continuous sequence in disk-to-
bulge ratio ("D/B"), which in turn
appears to be related to the specific
angular momentum (Fig. 3). Therefore,
since the frequency of morphological
types depends on the density of the
environment (Dressler 1980), the dis-
covery of two classes of E's is expected
to givenew insights in the relation be-
tween D/B ratios, the environment and
the process of acquiring angular
momentum during the primordial col-
lapse phase (see e. g. Barnes and Ef-
stathiou 1987). In addition, comparing
the disk properties along the D/B se-
quence is certainly important in order to
understand the mechanisms driving the
formation of these f1at components; this
is not a trivial aspect since disk proper-
ties (exponential decline and same cen-
tral surface brightness fto for all objects)
are suspected to be independent of the
D/B values (van der Kruit 1987,
Capaccioli and Vietri 1988).
Unlike disk-E's, whose structure
seems to result from initial conditions,
boxy E's are probably produced by re-
cent violent mergings. However, we
et al. (1988): - 1/4 of bright elliptical
galaxies exhibit nearly edge-on stellar
disks. From the statistics of inclination
angles, one therefore expects that up to
50 % of luminous elliptical galaxies
could in fact be disk-galaxies with very
low D/B ratios (D/B < 0.05).
Although the segregation of elliptical
galaxies into at least two distinct phys-
ical classes seems to be weil estab-
Iike classical 80's (the pointed
isophotes being reasonably interpreted
as caused by the superposition of an
edge-on stellar disk and an ellipsoidal
bulge). Actually, a strong morphological
similarity has already been firmly estab-
lished for those few "disk"-E's where
the faint edge-on disk has been disen-
tangled from the bright bulge (Carter
1987, Capaccioli 1987, Capaccioli et al.
1988). Although faint disks harboured in
disk-E's obviously have no direct dy-
namical consequence on host-galaxy
evolution, they are nevertheless symp-
tomatic of strong physical differences in
intrinsic shape and internal kinematics
as compared to boxy E's.
The physical nature of boxy E's is not
so clear. Boxy ellipticals are not only
anisotropic objects, but also present pe-
culiarities such as, e. g., kinematically
misaligned or even counter-rotating
cores (Franx and Illingworth 1988, Ben-
der 1988b, Jedrzejewski and Schechter
1988) and other features generally rec-
ognized as signatures of recent accre-
tion or merging processes (Nieto, 1988).
This strongly suggests that boxiness (at
least in these objects, see Nieto and
Bender, 1989) is related to these violent
phenomena. The distinct kinematical
differences between boxy and disky E's
shed light in particular on the large scat-
ter in velocity anisotropies of luminous
ellipticals (the (MT, (Vm/o)*) diagram,
Davies et al. 1983).
What fraction of elliptical galaxies be-
long to each class? A rough estimate
can be obtained from the investigation
of a sampie of 47 elliptical galaxies of
the Revised 8hapley-Ames Catalogue
of Bright Galaxies (8andage and
Tammann 1981) performed by Bender
( 0.50
cannot exclude that boxy E's just repre-
sent extreme examples of enivronmen-
tal evolution at early phases of galaxy
formation. The physical state of the hot
interstellar medium, which is present in
many of these objects, is certainly of
importance for the understanding of for-
mation and present-day evolution of
boxy ellipticals.
In summary, the isophotal shape anal-
ysis offers a new methodology/tool for a
complete revision of the classification of
early-type galaxies and our understand-
ing of the underlying formation mecha-
nisms/conditions, provided it is applied
to a significant and statistically mean-
ingful sampie of objects (in the same
way the Hubble classification scheme
was applied before). This revision obvi-
ously requires a huge research project.
Our key-programme is an attempt to
collect all the observational and theoret-
ical expertise to coordinate the efforts
toward an effective solution of the prob-
lem outlined above. As a first step of the
project, we aim to investigate the
occurrence of faint stellar disks in E
galaxies, their properties, origin, signifi-
cance, and relation to their analogues in
80's. 8urface photometry will be sup-
plemented by kinematical analyses in
order to investigate the kinematical be-
haviour of the spheroids and the dis-
tribution of their specific angular
momentum. Narrow-band images will
be used to study the properties of the
interstellar medium and relate the gas
properties to the other galaxian parame-
ters (shape of potential, X-ray emission,
Barnes, J., Efstathiou, G.: 1987, Astrophys. J.
Bender, R.: 1987, Ph. O. thesis, University of
Bender, R, Möllenhoff, C.: 1987, Astron. As-
trophys. 177,71.
Bender, R.: 1988a, Astron. Astrophys. Let-
ters 193, L7.
Bender, R: 1988b, Astron. Astrophys. Let-
ters 202, L5.
Bender, R, Oöbereiner, S., Möllenhoff, C.:
1988, Astron. Astrophys. Suppl. Series 74,
Bender, R., Surma, P., Oöbereiner, S.,
Möllenhoff, C., Madejski, R: 1989, Astron.
Astrophys. in press.
Bertola, F., Capaccioli, M.: 1975, Astrophys.
Binney, J.: 1978, Monthly Not. Roy. Astron.
Soc. 183,501.
Blumenthai, G. R, Faber, S. M., Primack,
J.R., Rees, M.J.: 1984, Nature 311, 517.
Burstein, 0., Oavies, R. L, Oressler, A., Faber,
S. M., Stone, R P. S., Lynden-Bell, 0.,
Wegner, G.: 1987, Astrophys. J. Suppl.
Series 64, 601.
Capaccioli, M.: 1987, in lAU Symp. 100:
Structure and Dynamics of Elliptical Galax-
ies, Ed. T. de Zeeuw, Reidel, Oordrecht, 47.
Capaccioli, M., Held, E. V., Nieto, J.-L.: 1987,
Astron. J. 94, 1519.
Capaccioli, M., Vietri, M.: 1988, in Yellow
Mountain Summer School on Relativistic
Astrophysics and on the Origin and Struc-
ture of Galaxies, preprint.
Carter, 0.: 1987, Astrophys. J. 312, 514.
Oavies, R L., Efstathiou, G., Fall, S. M., Illing-
worth, G., Schechter, P.: 1983, Astrophys.
Oavies, R L., Burstein, 0., Faber, S. M., Lyn-
den-Bell, 0., Terlevich, RJ., Wegner, G.:
1987, Astrophys. J. Suppl. Series 64, 581.
Ojorgovski, S., Oavis, M.: 1987, Astrophys. J.
Oressler, A.: 1980, Astrophys. J. 236, 351.
Oressler, A., Lynden-Bell, 0., Burstein, 0.,
Oavies, R L., Faber, S. M., Terlevich, RJ.,
Wegner, G.: 1987, Astrophys. J. 313, 42.
Fall, S.M.: 1983, in lAU Symp. 100: Internal
Kinematics and Dynamics of Galaxies, Ed.
E. Athanassoula, Reidel Oordrecht, p. 391.
Franx, M., lIIingworth, G.: 1988, Astrophys. J.
LeUers 327, L55.
Jedrzejewski, R.I.: 1987, Monthly Not. Roy.
Astron. Soc. 226, 747.
Kruit, P.C. van der: 1987, Astron. Astrophys.
Lauer, T. R.: 1985, Monthly Not. Roy. Astron.
Soc. 216, 429.
Nieto, J.-L., Bender, R.: 1989, Astron. Astro-
phys. in press.
Nieto, J.-L.: 1988, Invited Lectures at 2
Reunion de Astronomia Extragalactica,
Academia Nacional de Ciencias de Cordo-
ba, in press.
Nieto, J.L., Capaccioli, M., Held, E.V.: 1988,
Astron. Astrophys. Letters 195, L1.
Prugniel, Ph., Oavoust, E., Nieto, J.-L.: 1989,
Astron. Astorphys. in press.
Sandage, A., Tammann, G.A.: 1981, A Re-
vised Shapley-Ames Catalog of Bright
Galaxies, Carnegie Institution of Washing-
ton, Publication 635, Washington.
Silk, J.: 1983, Nature 301,574.
Gravitational Lensing
, U. BORGEESr
, o. LEFEVRE
C. VANDERRIESr
and G. WEIGELr
1 Institute ofAstrophysics, University of Liege, Belgium; 2 Canada-France-Hawaii Telescope Corporation, Hawaii;
3 Hamburg Observatory, F. R. Germany; 4 Astronomy Department, California Institute of Technology, Pasadena,
USA; 5 Physics Institute of the Erlangen-Nürnberg University, F. R. Germany; 6 Meudon Observatory, France;
7 European Southern Observatory, La Silla, Chile; 8 CITA, University of Toronto, Canada; 9 ESO, Garching, F. R.
Germany; 10 Haute-Provence Observatory, France
Prior to Professor van der Laan's en-
quiry, in the March 1988 issue of the
Messenger, on the general interest
among astronomers from the European
community to possibly participate in
Key Programmes (KPs) at the European
80uthern Observatory, at least three
distinct groups (including more than half
of the above authors) were already in-
volved in the study of "gravitational
lensing" effects (see box on pages
10-11). Observations were being per-
formed with the help of various tele-
scopes on La 8illa as weil as at other
observatories (VLA, CFHT, Palomar, Kitt
Peak, etc.).
A general feeling existed that our indi-
vidual work was progressing very slow-
Iy, the number of effective nights that
were allocated to our programmes be-
ing modest, also very much dependent
on unknown weather and/or seeing con-
ditions, on possible instrument failures
as weil as on the sometimes unpredict-
able decisions of the programme com-
mittees. We all knew that there was a
total absence of coordination between
our independent programmes and that,
of course, we could not avoid duplicat-
ing observations of similar objects. Fur-
thermore, each or our teams was hoping
very much to broaden its observational
interests: those studying "highly lumi-
nous quasars" wished to look also at
distant "radio galaxies", and vice-versa,
but this would remain a mere dream
until ... it became really possible for our
present team to submit an ESO Key
Programme entitled "Gravitational lens-
ing: quasars and radio galaxies". We do
feel very fortunate to-day since our pro-
gramme has been generously allocated
54 nights with the ESO 3.6-m and/or
NTI telescopes, 48 with the ESO/MPI
2.2-m telescope and 9 effective nights
with the 1.5-m Oanish telescope during
the next three years, starting effectively
during period 43 (April 1-0ctober 1,
1989). We should like to describe briefly
hereafter the main goals of our pro-
gramme and say a few words about the
different instruments and techniques
that we intend to use.
The Importance of Gravitational
Lensing Effects
Following the discovery of the first
example of a multiply lensed quasar
(Walsh et aL, 1979), there has been in-
creasing evidence that gravitational
lensing perturbs, to a significant extent,
our view of the distant Universe (see
Nottale, 1988 for a review on the sub-
ject). Gravitational mirages are indeed
being identified among QSOs, especial-
Iy highly luminous quasars (Sur-
deI et aL, 1987, 1988a-c; Magain et aL,
1988; and Ojorgovski, 1989),
among dlstant 3C and 4C radio galax-
ies (Le Fevre et aL, 1987, 1988a-b;
et aL, 1987; etc.) and as giant
lumlnous arcs lensed by individual
foreground galaxies and/or their associ-
ated cluster(s) (Soucail et aL, 1987a-b,
1988; Lynds and Petrosian, 1986, 1988;
Hammer et aL, 1988).
Highly significant statistical effects of
"gravitational amplification" have also
been reported for various sampies of
extragalactic objects. For instance, the
case of anomalous quintets - as weil as
other tight groups - of galaxies has
been accounted for by the lensing of
quartet haloes on background galaxies
(Hammer and Nottale, 1986a). Good
evidence for gravitational amplification
by intervening matter (stars, galaxies,
clusters, etc.) has been set forward for
the Brightest Cluster Galaxies (Hammer
and Nottale, 1986b), for quasars with
the richest absorption line spectra
(Nottale, 1987) and for a selected sam-
pie of flat-radio spectrum quasars (Fug-
mann, 1988). Furthermore, it has been
suggested that the variability of some
eruptive quasars such as 0846+513,
3C446, etc., is a direct consequence of
gravitational lensing by stars or com-
pact objects located in galaxy haloes
(Nottale, 1986). These so-called "micro-
lensing" effects had been predicted by
Chang and Refsdal (1979, 1984). Finally,
speculations have been made that the
second observed rise in the comoving
density of quasars from z '" 2.45 up to at
least z '" 3.8 (Veron, 1986) could be due
to statistical gravitational lensing effects
by foreground objects near the Iines-of-
A Quest for NewObservations
The answer to the question "what
fraction of extragalactic objects are
gravitationally lensed?" is very closely
related to that of "how do the visible and
dark matter distributions look at differ-
ent scales in the Universe?". It turns out
that any prediction made for the ex-
pected number of gravitationally lensed
quasars, radio galaxies, etc. is bound to
be very model dependent. We conclude
that it is essential to carry out a sys-
tematic observational search and study
of gravitational lensing effects in order
to better understand the luminosity
function of quasars, distant radio galax-
ies, etc., their observed number counts,
their apparent cosmic evolution and the
basic physical mechanism(s) powering
these energetic objects. The proposed
observations will also allow us to deter-
mine the Hubble parameter (Refsdal,
1964, 1966; Borgeest and Refsdal,
1984; Gorenstein et aL, 1988) as weil as
galaxy masses (Borgeest, 1986) on
account of the expected time delay be-
tween the brightness variations of multi-
ply lensed QSO images. More generally,
we shall obtain information on the dis-
tribution of luminous and dark matter at
various scales in the Universe. Further-
more, information on the size and struc-
ture of quasars should be derived from
the observation of micro-Iensing effects
(Grieger et aL, 1986, 1988).
Our Choice of Gravitational Lens
We do consider that the apparently
(mv< 18.5) and intrinsically (Mv< -29)
highly luminous quasars (hereafter
HLQs) as weil as the distant (z > 1)
powerful (P(178 MHz) > 10
radio sources (hereafter OPRSs) consti-
tute the best extragalactic candidates to
search for the presence of gravitational-
Iy lensed images at arcsec./sub-arcsec.
angular scale resolutions and/or for a
brightness amplification due to an ex-
cess of foreground objects (galaxies,
clusters) in the vicinity of the relevant
targets. The technical arguments lead-
ing to this assumption may be found in
Surdej et aL (1988c). High angular reso-
lution imaging of selected HLQs and
OPRSs with the ESO/MPI 2.2-m tele-
scope (plus direct CCO camera/OISCO
or digital speckle camera) under optimal
seeing conditions (FWHM < 1':2) is likely
to give important clues on our under-
standing of lensing effects by galaxies,
clusters and/or any other class of un-
known massive objects. Recent sys-
tematic searches for lensed QSO im-
ages with the 2.2-m telescope have al-
ready led to the discovery of three new
cases of multiply lensed HLQs: ESO
GL1 '" UM 673 (Fig. 2a), ESO GL2 '"
H1413+117 (Fig. 2b) and ESO GL3 '"
UM425 (Fig. 2c). Similarly, several mul-
tiple OPRSs have been identified among
the distant 3C radio sources at Mauna
Kea using the CFH telescope; strong
arguments supporting the mirage hy-
pothesis have been obtained for one of
them: 3C324 (Le Fevre et aL, 1987). Of
special interest is that a detailed com-
parison between CCO frames obtained
for a sampIe of 30 distant 3C radio
sources with z > 1 and randomly
selected fields indicates a significant ex-
cess of bright foreground galaxies and
Abell/Zwicky clusters near the 3C sour-
ces (Le Fevre and Hammer, 1989). It
seems that gravitational amplification by
foreground galaxies and rich clusters is
at least partly responsible fer the ob-
served radio and optical luminosities of
the bright 3C sources.
Further Proposed Observations
We intend to monitor under very good
seeing conditions (FWHM < 1':2) the
three ESO gravitationallens systems
mentioned above plus the Einstein
cross 2237+031 (Huchra et aL, 1985).
UM 673 (cf. Fig. 2a) constitutes our best
candidate to attempt an independent
determination of the Hubble constant
Ho, and hence to set an upper limit
to the age of the Universe, while
2237+031 appears to be the most ideal
object for detection of micro-Iensing
effects. It will be imperative to monitor
each of the four proposed targets about
once a week with the ESO 3.5-m ND
and/or ESO/MPI 2.2-m and/or 1.5-m
Oanish telescopes through a B filter. We
wish to thank in advance, for their com-
prehension, all observers at La Silla with
whom we shall routinely share observ-
ing nights during some 90 minutes.
Furthermore, because one expects
most of the multiply lensed QSO images
to have angular separations of less than
0':5 (e. g. Turner et aL, 1984), we wish to
image approximately 50 HLQs with
speckle masking. We recall that speckle
masking is an interferometric imaging
method that yields diffraction limited im-
ages with 0':05 resolution in spite of
image degradation by the atmosphere
Figure 2a: The double quasar ESO GL 1 = UM 673 (Nature 329
Figure 1b: Two different views of the VLA as seen in the early
morning of 17 January 1989 (Socorro, New Mexico).
Figure 2b: The Clover-Ieaf ESO GL 2 = H 1413+ 117 (Nature, 334,
Gravitational fields in the Universe may act on light rays emitted by distant sources in a way very similar to the refraction
properties of our atmosphere, or to the way lower air layers act on objects located near the horizon (cf. the Sun in Fig. 1a or
the Very Large Array or car lights at night in Figs. 1b-c).
Fig. 1d gives a schematic representation of the light ray paths when the ground turns out to be somewhat hotter than the
ambient air (cf. around noon on a sunny day). Because the refraction always bends light rays towards regions of colder air, a
second lower, inverted and somewhat deformed-image of the source (a palm-tree in this case) may result. Such a
multiplication, deformation and also amplification of different source images are readily seen in Figs. 1a-c.
Since the total solar eclipse of 1919, when astronomers observed for the first time an apparent displacement in the
positions of stars near the li mb of the Sun, it was recognized that light beams can be bent, not only in air layers having
different densities or in optical systems, but also in gravitational fields. As a matter of fact, this effect was predicted by
Einstein within his General Theory of Relativity. Bending of light is also observed when the light from a distant quasar,
and by telescope aberrations (Weigelt
and Wirnitzer, 1983).
In addition, since gravitational am-
plification of a compact or extended
source by large-scale foreground inho-
mogeneities (cf. rich galaxy clusters,
etc.) may very weil occur without multi-
plication of images, a detailed compari-
son of the count number of fjeld objects
(galaxies, clusters) will be made be-
tween the R GGD frames taken for the
HLQs, the DPRSs and randomly
selected fields. This unprecedented
statistical study combined with the ob-
served frequency of detecting multiply
lensed images versus their angular sep-
aration should enable us to quantify the
importance of lensing effects as a
We also plan of course to perform
detailed photometric and spectros-
copic studies of the immediate
surroundings and projected intergalac-
tic medium near known, suspected and
expected new ESO gravitationallens
Our recognition naturally goes to Prof.
van der Laan for having promoted so
convincingly the idea of the ESO Key
Programmes in the European communi-
ty. Our thanks are also due to the Ob-
serving Programmes Gommittee and to
the external referees for their patient
(and somewhat hidden) work. We spe-
cially thank Jacques Breysacher and
Ghrista Euler, for gently orchestrating
the non trivial task of scheduling all OUr
Figure 1d: Explanatory diagram: formation of terrestrial mirages.
hot air
cold air
refraetion of-
light rays
Figure 1c: Lights from a distant car on the national road between
Magdalena and Gatil (New Mexico, 16 January 1989).
"0 • 1Wl• .r;.. .,r' " I" ' .. ....'il
·.0.,.. ,..' tIl•• d t-. _, •
- fI' .. •• _",...,
" ". •.
. , ...' - ".... .........
. , r • .. iI "r· ...
.' • 0 .. 1 oe' ,. 41.. l!1. • l
"". • • • &..J iII .. .' ..
• ••v . .;(J".••• '•. ..,.. ,'''i\9''
..... - • .. '. '.1' ...
• . ....... : '1 • • ,. .; .:. .....
. "- ... - •... . f""'\.-II.... E:
· 0.....:.... ..
.•••••••••. '"
., '. ... " . . ....... ...
. •• " _. r·.. ••• , .. _. • •
: .... . - .. .
gravitational raY2
Figure 2c: The multiple quasar ESO GL 3 =UM 425 (Ap. J. Letters,
in press).
Figure 2d: Explanatory diagram: formation of gravitational mirages.
galaxy or any other astronomical source passes close by one or more massive objects on its way to uso Such objects
may be individual galaxies (cf. Fig. 2d). clusters of galaxies or even larger structures in the Universe. This effect is
referred to as the so-called "gravitational lensing". Depending on the intensity and form of the gravitational field, the light
from the quasar may not only be bent into multiple images, but some of these images may become brigllter than the
quasar itself would have appeared in the absence of the gravitationallens (cf. Figs. 2a-c). This is referred to as "light
amplification". Due to the amplification effect, we may be able to observe gravitationally lensed images of very distant
quasars, galaxies, etc. which would otherwise have been too faint to be detected with present teles·copes. Gravitational
lenses may therefore act as giant telescopes, allowing us to investigate otherwise inaccessible, very remote regions of
the Universe. It is partly in order to evaluate the extent to which our view of the distant Universe corresponds to a still
unveiled mirage, and not to the real Universe, that we have proposed to perform the studies summarized in the present
requested observations.
Finally, we apologize for not having
quoted in this short note the works of all
scientists having contributed to the pro-
gress of our knowledge in the field of
gravitationallensing; they are simply just
too numerous!
BorgeeSl, U.: 1986, Ap. J. 309, 467.
BorgeeSl, U., Refsdal, S.: 1984, A & A 141, 318.
Chang, K., Refsdal, S.: 1979, Nature 282, 561.
Chang, K., Refsdal, S.: 1984, A & A 132,168.
Fugmann, W.: 1988, A & A 204,73.
Gorenslein, M. V., Falco, E. E., Shapiro, 1.1.: 1988,
Ap. J. 327, 693.
Grieger, B.. Kayser, R., Refsdal, S.: 1986, Nature
Grieger, B., Kayser, R.. Refsdal, S.: 1988, A & A
Hammer, F., Notlale, L.: 1986a. A & A 155, 420.
Hammer, F.. Notlal , L.: 1986b, A & A 167, 1.
Hammer, F., Le Fevre, 0., Jon s, J., Rigaut, F.,
Soucail, G.: 1989, A & A 208, L7.
Hammer, F., Le Fevre, 0., Notlale, L.: 1987, in
Observational Cosmo/ogy: Proceedings or lAU
Symposium no. 124, ed. by Hewill et al., Reidel,
Dordrecht, 751.
Huchra, J., Gorenslein, M., Kenl, S., Shapiro, 1.,
Smith, G.. Horine, E.. Perley, R.: 1985, As/ron. J.
Le Fevre, 0., Hammer, F., Notlale, L.. Mathez, G.:
1987, Na/ure 326,268.
Le Fevre, 0., Hammer, F., Notlale, L., Mazure. A.,
Christian, C.: 1988a, Ap. J. 324, L 1.
Le Fevre, 0., Hammer, F.. Jones, J.: 1988b, Ap. J.
Le Fevre, 0., Hammer, F.: 1989, in preparation.
Lynds, R., Pelrosian, V.: 1986, Bull. Amer. As/ron.
Soc. 18, 1014.
Lynds, R., P lrosian, V.: 1988, Preprint.
Magain, P., Surdej, J., Swings, J.P., Borgeesl, U.,
Kayser, R., Kuhr, H., Refsdal, S., Remy, M.: 1988,
Nature 334, 325.
Meylan, G., Djorgovski, S.: 1989, Ap. J. (Letlers), in
Nottale, L.: 1986, A & A 157, 383.
Nottale, L.: 1987, Annales de Physique 12, 241.
Nottale, L.: 1988, Annales de Physique 13, 223.
Refsdal, S.: 1964, MNRAS 128,307.
Refsdal, S.: 1966, MNRAS 132,101.
Soucail, G., Mellier, Y., Fort, B., Picat, J.P.: 1987a,
Soucail, G., Mellier, Y., Fort, B., Hammer, F.,
Mathez, G.: 1987b, A & A 184, L7.
Soucail, G., Mellier, Y., Fort, B., Mathez, G.,
Cailloux, M.: 1988, A & A 191, L19.
Surdej, J., Magain, P., Swings, J.P., Borgeest, U.,
Courvoisier, T.J.-L., Kayser, R., Kellermann, K.I.,
KUhr, H., Refsdal, S.: 1987, Nature 329,695.
Surdej, J., Magain, P., Swings, J.P., Borgeest, U.,
Courvoisier, T.J.-L., Kayser, R., Kellermann, K.I.,
Kuhr, H., Refsdal, S.: 1988a, A & A 198, 49.
Surdej, J., Swings, J.P., Magain, P., Borgeest, U.,
Kayser, R., Refsdal, S., Courvoisier, T.J.-L.,
Kellermann, K.I., Kuhr, H.: 1988b, Publ. Astron.
Soc. Pac. Suppl., in press, ESO preprint no. 571
(Proc. NOAO Workshop on Optical Surveys for
Surdej, J., Magain, P., Swings, J.P., Remy, M.,
Borgeest, U., Kayser, R., Refsdal, S., Kuhr, H.:
1988c, in the proceedings of the First DAE.C.
Workshop "Large scale structures: observations
and instrumentation", Paris, eds. C. Balkowski
and S. Gordon.
Turner, E.L., Ostriker, J.P., Gott 111, J.R.: 1984,
Ap. J. 284, 1.
Veron, P.: 1986, A &A 170, 37.
Walsh, D., Carswell, R.F., Weyrnann, R.J.: 1979,
Nature 279, 381.
Weigelt, G., Wirnitzer, S.: 1983, Opt. Lett. 8, 389.
The New Research Student Programme of the European
Southern Observatory
H. VAN DERLAAN, Director General, ESO
For many years ESO has appointed
young astronomers as Fellows to work
one to two years at Headquarters or on
La Silla. The HO Fellows spend the
greater part of their eftorts on personal
research, the Fellows on La Silla split
their time roughly 50/50 between their
research and support astronomy duties.
The ESO Fellowship programme has
successfully contributed to the develop-
ment of young scientists from member
states into mature research astronom-
ers. It has also promoted the interaction
among investigators from many piaces
and traditions, interactions that continue
long after the collaborators have left
their ESO posts. This postdoctoral
Fellowship programme continues undi-
Henceforth ESO is also oftering Stu-
dentships, about 8 appointments per
year, equally shared between Garehing
and La Silla. This article introduces this
new predoctoral Studentship Pro-
Selection and Studentship
ESO aims at having sixteen such
graduate students in the Organization's
establishments at any time. The
programme will be semi-annually
announced in the Messenger and by the
circulation of announcements in
member states' institutes. Applicants
have to make use of a form designed for
the purpose and available upon request
from Personnel and General Services at
Headquarters. Deadlines for application
will be May 1st and December 1st.
Appointments can commence through-
out the year.
Students whose duty station is the
Garching Headquarters, will be required
to spend at most 25 % of their official
working time on duties in support of
Visiting Astronomers (users of measur-
ing machines, of computers with the
IHAP and MIDAS data reduction sys-
tems, and astronomers using the Re-
mote Control Observing Facilities). Stu-
dents on La Silla will normally be part of
the team in the La Silla astronomy de-
partment that provides introductions
and observational support to Visiting
Astronomers. Typically, half their ofticial
working time will be occupied by these
Students are appointed for one year,
normally renewable by a second, final
year. During their studentship term, they
may have diverse forms of support from
their home institution. ESO will supple-
ment these to a maximum, total stipend
of DM 2200/month during the first year.
For many years young astronomers
and engineers from France, called
cooperants, have worked on La Silla in
national service, which is an alternative
to military duty. More recently Belgium
and Italy also ofter such "cooperant"
programmes. Henceforth ESO will re-
quire cooperants, whether at HO or on
La Silla, to meet at least the same re-
quirements as candidates for student-
ships. Cooperants are normally pre-
selected by national selection com-
mittees and then proposed to ESO.
Some cooperants already have a doc-
tor's degree and are then regarded as
Fellows. Where they are selected by
studentship criteria the cooperants will
henceforth be counted as part of the
contingent of 16 students.
Selection Criteria - Necessary
- Candidates must be registered post-
graduate students at a recognized
university in an ESO member state.
- Candidates must have at least an out-
line of their Ph. D. dissertation pro-
gramme and a professor who accepts
responsibility for that programme and
the student's supervision.
- The candidate must have fulfilled the
normal course- and examination re-
quirements for the Ph. D., except the
dissertation research.
- The candidate's programme must be
such that it can be successfully pur-
sued at ESO HO or on La Silla, be-
cause of its aftinity with work going
on there and facilities available for the
- The candidate's university must
guarantee the conditions necessary
for the student to complete her/his
dissertation after the tenure of the
ESO Studentship.
- An ESO staft member must be pre-
pared to be the student's local super-
visor and mentor, notwithstanding the
continuing responsibility of the stu-
dent's university supervisor.
From this list it should be obvious that
ESO in no way seeks to play the role of a
degree-granting institution. The pro-
gramme means to provide oppor-
tunities, circumstances and facilities
which enhance the participating univer-
sities' postgraduate programmes and
enrich selected students' early research
experience. The conditions listed above
are necessary. From among applicants
potentially capable of meeting them,
ESO will select those whose talents and
circumstances are best attuned to the
programme's goals.
Groups Where Research Students
There is a brochure in preparation,
available in a few months, which will
describe the programme in some detail.
For the moment the following brief indi-
cations suffice. It is expected that most
of the studentships will be offered in the
Astronomy Group of the Science Divi-
sion at Headquarters in Garching and in
the Astronomy Department at the La
Silla Observatory. In addition, more in-
strumentally oriented students may
work in one of three groups in the Prp-
ject Division at Headquarters, viz. Opti-
cal Instrumentation, Infrared Instrumen-
tation and High Resolution Imaging and
Interferometry. Computer science stu-
dents may find possibilities in the Image
Processing Group of the Science Divi-
sion, and occasionally physics or en-
gineering students can be engaged in
one of the three discipline groups of the
Technology Division, viz. Electronics,
Optics and Mechanics.
On La Silla, the latter interests may
also be served in the department called
Technical Research Support. The just
mentioned technology disciplines are
also practised in the TRS. Finally the
SEST group, running the Swedish-ESO
sUbmillimetre telescope on La Silla, and
the group known as VLT Site Services
provide a good context for a variety of
astronomical, technical and atmospher-
ic research interests.
ESO Motives for the Studentship
There are many, most quite obvious,
and I mention but a few. ESO is firstly a
service organization for European as-
tronomy. To fulfil this mission there is a
heavy operational, technological em-
phasis in what we do: the reliability of
the service is paramount. All this activity
aims at enabling the best astronomical
research attainable. The quality of this
work requires an acute awareness of
astronomy's requirements. Having re-
search students in our teams in addition
to fellows, to visitors and to staff serves
this purpose: research students, bright,
ambitious, naively demanding, contri-
bute to the research mindedness of an
institution. This needs to be maintained
both in Garching and on the mountain.
Secondly, the studentship pro-
gramme serves to extend the Iinkage of
ESO to Europe's universities, of ESO
staff to the academic staff of the as-
tronomy institutes in its users communi-
ty. Such linkage is indispensible for
the continuous information flow that
keeps our priorities and services attuned
to the research requirements of the
A third aspect concerns the long-term
quality and ambitions of European as-
tronomy embodied in the next genera-
tion. The fellows and students who
spend a year or two within ESO are
better equipped to use its facilities for
their personal research in future. In addi-
tion, and just as important, they will
enable the institutes that employ them
to use ESO telescopes and services to
best advantage. They are the vanguard
of VLTobservers, training now and
set to work for decades in the next
Finally, and related to the preceding
point, these youthful scientists establish
patterns of professional and personal
relations among themselves, relations
that will guide their collaborations and
projects of the future. The result will be a
lowering of national boundaries, the
growth of European excellence in as-
tronomy. And that too is ESO's raison
The Users Committee (UC)
B. MARANO, UC Chairman, Osservatorio Astronomico, Bologna, Italy
The Committee meets at least once a year. It is convoked by the Director General.
The members (one from each member country) are appointed by the Director General
from among the recent Visiting Astronomers for four year terms (not immediately
renewable). The terms are staggered, so that each year two persons are replaced. The
Committee annually selects its Chairman. National Committees are invited to submit
nominations for membership to the Director General.
heard two reports on key programmes
and ESO policy on data archiving.
The atmosphere is always sound and
collaborative, and plenty of time is de-
voted to the discussion of each item.
The effectiveness of the Users Com-
mittee must also be evaluated on
B. Marano (Chairman), Italy, 1986-1989
B. Stenholm, Sweden, 1989-1992
J. Surdej, Belgium, 1987-1990
Chr. Trefzger, Switzerland, 1986-1989
Present Composition of the UC
M. Azzopardi, France, 1988-1991
L. Hansen, Denmark, 1987-1990
F. P. Israel, Netherlands, 1986-1989
J. Krautter, F.R. Germany, 1989-1992
The Committee advises the Director General on matters pertaining to the functioning
of the La Silla Observatory from the point of view of the Visiting Astronomers. It should
consider the possibility to arrange a Users Conference.
scopes on La Silla, deterioration of the
dome seeing at the 3.6-m telescope,
flight safety.
The agenda is often closed by some
specific items, related to current or
foreseen projects. In the last meeting,
for instance, the Users Committee
The Users Committee meets as a rule
once a year, in May. Besides the nation-
al representatives, the meeting is
attended by the Director General and by
the ~ e a d s of La Silla, of the Image Pro-
cesslng and Measuring Machines and
of the Visiting Astronomers Office. Other
memb.ers ?f the ESO staff may attend,
especlally If they are directly involved in
some item in the agenda.
The meeting is opened by areport of
the Director General on the present ESO
state and future perspectives. The Head
of La Silla then reports on more techni-
cal subjects, as, for instance statistics
of failures at the telescopes ~ n d recent
or foreseen improvements in the in-
strumentation. A similar report on image
processing and measuring machines of-
ten folIows.
National representatives are then
called to speak. They have collected
from their colleagues complaints, im-
pressions, suggestions. Their reports
can span from the description of prob-
lems pointed out by a single observer or
group to very general matters. A few
examples of general issues raised in the
last two meetings can better give the
spice of the discussion: introduction of
service observing, future of small tele-
All meetings will take plaee at ESO in
back to a not-very-interesting "safety
valve for disgruntled astronomers". The
choice is mostly up to uso
Tentative Time-table of
Council Sessions and
Committee Meetings in
another ground: does action follow the
discussion? My experience is that,
when specific technical problems are
identified, they are solved by ESO in
relatively short time. When a long-term
action or a change in ESO policy are
required, or when budget problems are
involved, the Users Committee repre-
sents only one of a number of steps in
the process, and things are obviously
not so simple. In my opinion it is widely
feit that the role of the U. C. in these
circumstances could be better clarified.
In the past, some efforts have been
made to improve the work of the U. C.
and its effectiveness. It has become
customary to have an informal gathering
of the national representatives the day
before the annual meeting. The issues
raised by various members can be com-
pared, and common and general prob-
lems can be more easily extracted and
presented in the meeting after deeper
consideration. Furthermore, it has been
realized that pointing out a general
problem in a wide, multinational com-
munity is often a slow process. Solving
the problem can take a long time as
weil. The overall process can barely be
followed if its typical timescale is longer
than the turnover time of the members
of the Committee. For this reason it was
proposed, and approved by the Director
General and by the Council, to extend to
four years, that is to four meetings, the
term of the members of the U. C.
Looking to the near future, one can
foresee several changes in our way of
working at the telescopes: key pro-
grammes imply a different way of
scheduling and using them; the availa-
bility of both the 3.6-m and the ND will
permit more flexibility in the instrumen-
tation; remote observing is becoming a
real possibility; flexible scheduling is
currently proposed, in various obser-
vatories, as a way of better exploiting
optimum sky conditions. These exam-
pies, only a few from a longer list, show
a strongly evolving situation. The users
can play a critical role in it, providing
essential inputs and acting as a feed-
back. The Users Committee could be an
important link in this process. Or, in
absence of a continuous pressure from
the community of users, it could slip
May 10-11
May 18-19
May 30-31
Nov. 30-Dee. 1
Users Committee
Finanee Committee
Seientifie Teehnieal
Observing Programmes
Committee of Couneil
Seientifie Teehnieal
Finanee Committee
Observing Programmes
Committee of Couneil
A. BLAAUW, Kapteyn Laboratory, Groningen, the Netherlands
"- - - observers are on duty fram sunset till sunrise - - -. "
From Andre Muller's instruetions for the site tests, Deeember 1960.
Over a time span of more than seven
years, with several interruptions from
late 1955 to the middle of 1963, young
European astronomers and their assis-
tants have been engaged in the search
for a site in South Africa. By the end of
that time, it became clear that the ob-
servatory would not be built on this con-
tinent; the South American Andes
Mountains offered superior observing
Does it make sense, then, to devote a
full chapter to the South African explora-
tions? It does - not only because we
want to do justice to the large effort
made by many young astronomers and
their assistants, but also because the
South African venture was ESO's first
exercise in European collaboration.
First Impressions
Already in January 1954, at the sec-
ond meeting of the ESO Committee
(henceforth to be denoted by EG), the
• Article No. I appeared in (he Messenger of Oe-
cember 1988.
question of the best site for the observa-
tory was taken up. As I explained in the
previous article, the southern part of
Africa seemed a natural choice. How-
ever, the major observatories in South
Africa were all located in, or near, major
cities er communities: the Cape Obser-
vatory, the Union Observatory - origi-
nally only at Johannesburg but later
having its field station at nearby Har-
tebeespoortdam -, the Boyden Obser-
vatory near Bloemfontein, and the
Radcliffe Observatory near Pretoria.
This latter observatory had been created
rather recently, in the early 1930's, as a
result of the transfer of facilities from
Oxford; yet also in this case proximity to
a major city had been chosen, even for
the planned 74-inch telescope [1].
For ESO, vicinity of a major centre of
civilization was not an important crite-
rion, and so, the EC decided to start
from scratch. Needed was, of course, a
place with a minimum of cloudiness and
as free as possible from smoke and sky
illumination. Moreover, astronomers
want good "seeing". By this they mean,
that the image of a star as observed in a
telescope should show minimum distor-
tion due to turbulence in the earth's
atmosphere. This question of "seeing" is
explained in some more detail in the box
accompanying this article.
Apart from the experience collected
over the years by the existing obser-
vatories, there was little the EC could go
by. There was an interesting report by
B. J. Bok of August 1953, dealing with a
comparison of conditions at Harvard
Observatory's Boyden Station in South
Africa and its Agassiz Station in
Massachusetts [2], in which Bok drew
attention to what seemed to be a gener-
al characteristic: "All over the High Veld
of South-Africa, with its remarkably
clear and pure skies, the seeing de-
teriorates often about midnight or short-
Iy after, with no recovery before
dawn - - -. The after-midnight deterio-
ration of seeing happens as weil at the
Union Observatory in Johannesburg, at
the Radcliffe Observatory near Pretoria
Astronomical "Seeing"
Under ideal atmospheric conditions, the image of a star as seen in a telescope consists
of a bright central spot surrounded by a weak circle, the diffraction ring. This is due to the
wave character of the light in combination with the fact that the telescope objective or
mirror cannot but be of limited size. If the atmosphere is disturbed by turbulence, then (a)
the ring is broken up and both it and the central spot loose their sharpness, and (b) the
whole image of the star moves rapidly, in an erratic manner. The combination of these two
effects determines the quality of the image, called seeing by the observer. The less
turbulent the atmosphere, the better is the seeing. Hence, astronomers can judge seeing
by the quality of the appearance of the diffraction ring, and by the degree of violence of the
motion of the bright central spot, called the "image motion".
Estimates by the appearance of the diffraction ring are not easy to put on a quantitative
basis; observers use ascale of ratings mutually agreed upon and to be checked regularly.
The Danjon telescopes, equipped with mirrors of 25 cm, produced a suitable size of the
diffraction ring and gathered sufficient light to make it weil visible for bright stars. Normally,
the estimates were not seriously hampered by the image motion.
Judging the seeing by the image motion has the advantage of allowing a quantitative
measurement, for instance the average deviation of the central spot from its mean position.
It has the disadvantage of requiring very stable mounting of the telescope. There is,
however, a way around this: one fixes two telescopes on one sturdy mounting and by
means of an optical device arranges for the two fields of view to be seen superposed on
each other. Measurement of the relative displacement of the two central spots is then a
measure of only the atmospheric effect because the shaking of the mounting affects the
two in the same manner.
There is still a third method that helps measuring the seeing. In a turbulent atmosphere,
we can distinguish turbulence cells, somewhat vaguely defined units which move with
respect to the surrounding medium. Such cells differ slightly in temperature with respect to
this medium. As a consequence, if one measures the temperature at a fixed point above
ground level - for instance at the top of a fixed pole - then one will find rapid fluctuations as
a consequence of the successive passages of the cells and the surrounding medium. The
more turbulent the atmosphere, the more violent the temperature fluctuations. Experiments
have shown that the degree of violence is closely correlated to the rating of the seeing by
the diffraction ring method or by the image motion. Conversely, measures of the tempera-
ture fluctuations can tell us whether we may expect to observe with good, or with bad
and at the Lamont-Hussey Observatory
on Naval Hili in Bloemfontein. - - -."
Also of historical interest is an exten-
sive letter by Walter Baade to Oort of
1 November 1954 [3]. That Baade's
opinion would carry much weight is ob-
vious: his fundamental discovery of the
different stellar populations had been
Possible by a combination of two spe-
cial circumstances at Mt Wilson Obser-
vatory some time during World War 11: a
sky free of illumination by the neigh-
bouring city of Pasadena, and excep-
tional seeing conditions at the 100-inch
during the photographic exposures of
the Andromeda Nebula and its satellites.
Baade's letter stressed the impor-
tance of local conditions: "- - - I have
no experience with the conditions on
high plateaus such as that in South Afri-
ca but I am strongly inclined to believe
that there, just as in Southern California,
the seeing during the best observing
season is largely determined by the air
layers close to the ground - - -. Local
topological conditions therefore must
playa role - - -." Baade also stressed
the importance of correlating the rating
of the seeing as judged in the test instru-
ment with that observed in a large re-
flector, and suggested that the Haute-
Provence Observatory, favoured with
good seeing, might be a suitable place
for such comparison.
At the request of the EC, meteorologi-
cal data on South Africa were collected
and discussed by Siedentopf of Tübin-
gen [4]. He concluded that the High
Veld, the semi-desert plateau stretching
from Johannesburg to Bloemfontein
and southward, should offer the most
favourable over-all conditions. The EC
m e ~ t i n g of November 1954 therefore
declded to first explore the Pretoria-
Johannesburg and Bloemfontein-Kim-
berley areas, with limited tests in the
Beaufort West region located further
south. In each of the first two a fixed
observing post was to be chosen near
the existing observatory, to serve as a
reference point, and the surroundings
were to be explored with a moving tele-
At a meeting in March 1955 in Uccle,
details of the project were discussed by
Bourgeois, Danjon, Heckmann, Spencer
Jones and Oort [5]. A classical method
was chosen for the evaluation of the
quality of stellar images: the appearance
of the diffraction rings as observed in a
small reflector. In the accompanying
box we explain some of the ways in
which the astronomer can evaluate the
quality of the stellar images. The method
selected had been described by Danjon
and Couder in their textbook Lunettes et
Telescopes of 1935 [6]. Four azimuthally
mounted reflectors of 25 cm aperture
were built for the praject at the Paris
Observatory. For the measurement of
the atmospheric extinction, photo-elec-
tric observations were to be made at
wavelengths about 4500 and 5300 A,
with small refractors. Moreover, of
course, cloudiness, wind velocity and
wind direction would have to be re-
The first observers lett in October
1955 by boat: G. Courtes from France,
J. Dommanget from Belgium, H. Elsäs-
ser from the German Federal Republic,
and Ch. E. Heynekamp fram the Nether-
lands. They arrived in Capetown on
November 6. An extensive letter by El-
sässer and Heynekamp to the spiritual
father of the project, J. H. Oort, of 17
January 1956, reports on the beginning
of their activities [7]. Elsässer described
this early work in Die Sterne 33, p. 3,
1957. First observations for intercom-
parison of all observers were made in
the Bloemfontein area in December
1955. Subsequently, the work was di-
vided over the northern (Johannesburg-
Pretoria) area and the southern one
On the basis of this first reconnais-
sance, the EC meeting of April 1956
decided to drop the Johannesburg-Pre-
toria area and concentrate further work
on the region of Bloemfontein and farth-
er southward down to the surroundings
of the village Oudshoorn. This is located
close to the Swartberg mountain range
at the southern border of the Karroo
semi-desert. J. Boulon from France re-
placed G. Courtes for this second
phase, which was reported at the Oc-
tober 1956 meeting of the EC by Danjon
and Siedentopf. The Karroo emerged
more and more as the most promising
region, so that the EC decided to extend
testing there, in particular near the set-
tlements Zeekoegat and De Aar-
Koffiefontein. A new team of observers
replaced those mentioned before: F.
Bertiau from Belgium, and K. Rohlfs and
J. W. Tripp fram the FRG. They em-
barked upon a years' programme to be
completed by September 1957. A joint
interim report of March 23, 1957 submit-
ted to the chairman of the EC [8] gives
first observational results, but also re-
flects some concern about the problem
of systematic differences between the
results of different observers, the rele-
vance of tests made with small tele-
scopes for work done with large tele-
scopes, the possible influence of local
seeing-disturbing elements, etc.
Nevertheless, the Zeekoegat area be-
gan to seem superior to the other ones
After the completion of this mission in
September 1957, the work done over
the years 1955-1957 was reviewed in a
meeting on 9-11 January 1958 at the
Paris Observatory by a group consisting
of Danjon, Heckmann, Fehrenbach,
Gouder, Dommanget, Guinot and Tripp
[9], which led to are-analysis of the data
by Tripp [10]. This was first prepared for
the July 1958 meeting of the EG and, in
more complete form, for its October/
November meeting. It confirmed the
favourable impression of the Zeekoegat
site, with Tafelkopje, a hili near Bloem-
fontein, as a close second. However, as
the report pointed out, the analysis
suffered from systematic differences
between the evaluations obtained by
different observers at different places
and between observations made with
different telescopes, notwithstanding
the careful measures taken to eliminate
these effects. An independent analysis
was published by Dommanget [11].
Adding Some "Real" Astronomy
A new phase of more rigorous investi-
gation in the Karroo developed in the
second half of 1958. It envisaged, apart
from continuation of work with the Dan-
jon and photo-electric telescopes, some
real research programmes. A sugges-
tion for such broadening had been
made by Danjon at the July 1958 meet-
ing of the EG. Doing "real" astronomy
would help testing the site and make it
more attractive for young astronomers
to become involved in the work. Two
projects presented themselves for this
The Marseilles Observatory, directed
by Fehrenbach, had developed plans for
the erection in South Africa of a dupli-
cate of the objective-prism refractor of
40 cm aperture of the HPO, the so-
called GPO (Grand Prism Objective) for
determining radial velocities [12]. Initial-
Iy, a location near the village Prince Al-
bert had been considered, but now an
alternative was contemplated: one of
the possible sites for the ESO project.
The operation would become more ex-
pensive for logistic reasons; the addi-
tional costs might then be absorbed by
the site testing project [13]. The solution
also would strengthen the effort for in-
tegration of France in the ESO project.
At this time the French participation in
ESO was still quite uncertain.
A second suggestion had been made
by myself on behalf of the Kapteyn
Laboratory: it proposed photo-electric
't0::! .,.....,
TAuLt MOUNTAIN
on K/lwtlt'v/.i Sft. 10 KI.v.rvJ.1
___-:::---::-:-:----"" '.,m
Gcneu O
\ .
Prelom North
Map of South Afriea. The ESO site tests over the years 1955-1963 eovered the area from around Pretoria-Johannesburg to the Great Karroo
near Ihe south eoast, and during the last years eoneentrated on the region marked with the rectangle. This region is shown in blow-up in IM
upper left corner, adapted from the report by Ursula Mayer mentioned in Ihe text. The sile testing station near Zeekoegal and Ihe Ihree on
Klavervlei Farm: Table MI., Roekdale MI. and Flathi/l are marked by crosses. The blow-up of Table Mt. on Klavervlei Site shows the Ihree
loeations investigated by the Quick-Look expedition in early 1961.
The Zeekoegat Site.
Above: In spring 1959: members of the "Technical Group" with the
Owner of Sunny Side Farm; from left to right, Fehrenbach, Haffner,
MIss Oosthuizen, Hooghoudt. Photograph by the author.
Below: ESO's "Zeekoegat Station" in 1962 in the background ESO
buildings and houses. Photograph by 0.
The Klavervlei Site.
Above: Table Mountain (middle, background) and in front of it the
Klavervlei Farm Settlement. Photograph by the author, 1959.
Below: Klavervlei Farm Settlement seen from the air; the large dark-
roofed house in ihe lower left housed the ESO observers. Photo-
graph by D. Beintema, 1962.
Photometry of moderately faint stars,
Information on the photomet-
rlC quallty of the site as a by-producL As
It turned out later, this project could not
be realized, but a similar one was done
by the Tübingen programme of Sieden-
topf described below.
These suggestions were submitted to
the October/November 1958 meeting of
the EC [14]. at which also another step
was taken up: an evaluation of building
costs, technical expertise, acquisition of
water and power supply, etc. in South
Africa. These were to be investigated by
a "technical" group consisting of the
englneer B. G. Hooghoudt (responsible
for technical developments of radio as-
tronomy in the Netherlands), Fehren-
bach and myself, together with the Ger-
man astronomer H. Haffner who at that
time resided at the Boyden Observatory.
The group arrived in South Africa on 16
March 1959 and stayed for about five
weeks, after which it reported to the EC
meeting of May 1959 [15].
The report led to a somewhat mod-
ified approach. Further testing of the
Zeekoegat area was recommended, but
attention was also to be given to sites at
Considerably higher elevation than those
explored so far. Such sites were to be
found on the Nieuwveld Plateau north-
west of the village Beaufort West. On
the other hand, no further testing of the
region around Bloemfontein was to be
done. Reasons for its exclusion were the
fear for growing disturbance by city
lights, and seasonal effects in the cli-
mate which are unfavourable for observ-
ing tl1e Magellanic Clouds.
Henceforth, interest focussed mainly
on two possible locations: the vast
territory of Klavervlei Farm on the
Nieuwveld Plateau, where contacts with
the owner, R. Köster, had been estab-
lished by the "technical" group; and
Zeekoegat, where the same had been
done with the owner Miss M.E.Z. Oost-
huizen of the Farm Sunnyside. Klavervlei
Farm was located about 35 km north-
west of Beaufort West, and Zeekoegat
about 80 km south of this town. Ulti-
mately, three montain spots on Klaver-
vlei Farm became the subject of inten-
sive tests: Table Mountain at elevation
about 1,970 m, Rockdale Mt. at
1,860 m, and Flathill at 1,490 m. They
are indicated on the accompanying
The Zeekoegat site was located at
elevation about 1,000 m, only slightly
above the surrounding plane. In a way,
the two kinds of sites represented two
different philosophies: in the Klavervlei
area, the mountain-top concept em-
bodied by the Californian observatories;
at Zeekoegat the concept of the French
Haute-Provence Observatory - only
slightly elevated above its surroundings
- which reminds us of the description at
the end of the chapter on image quality
in Danjon and Couder's Lunettes et Te-
lescopes referred to before: "O'une
maniere generale, il convient de re-
chercher de preference les plateaux
secs d'altitude moyenne, loin de la mer
ou des grandes va/lees, couloirs de
vent. /I est superflu d'avoir un horizon
degage, car un observatoire astronomi-
que n'est pas un point de vue. - - _n.
The Quick-Look Expedition
In order to get a first impression of the
Nieuwveld Plateau, one of the Klavervlei
sites, Table Mt. was explored by a
three-month "Quick-Look" expedition.
However, whereas the earlier tests had
been limited to weil accessible loca-
tions, for the Klavervlei sites road con-
struction was a first requirement. It was
achieved through the intermediary of the
owner of the farm, so that in September
1959 access to Table Mt. was possible.
11 km of roads suitable for four-wheel-
drive vehicles were constructed, leading
to the three observing locations num-
bered I, 11, and 111 on Table Mt., marked
on the accompanying map.
The Quick-Look expedition was
carried out by Andre Muller in collabora-
ti on with the Swedish geodesist C. Ulff.
Andre Muller was one of my associates
at the Kapteyn Laboratory and had pre-
viously conducted observations at the
Leiden Observatory Station on the pre-
mises of the Union Observatory at
Johannesburg. He was, therefore, weil
acquainted with South Africa. In
November 1959, on the way to South
Africa, Muller and Ulff spent a week at
the Haute-Provence Observatory in or-
der to gain experience with the use of a
Danjon telescope in consultation with
the staff of the HPO. They completed
the Quick Look per 1 April 1960 after
three months of seeing tests and c1i-
mate monitoring, and Muller reported at
the July 1960 meeting of the EC [16].
Letters of Muller to myself in the
period December 1959 to March 1960
[17] describe delays in the transport of
the telescopes and the rather primitive
living conditions under which the Quick
Look had to be executed (for shelter
during the night a tent was borrowed
from the Dutch organization ZWO), and
troubles with the instruments, among
which a lack of stability of the mounting
of the Danjon telescopes under the
sometimes very strong (and cold!) winds
on Table Mt. For the follow-up of the
Quick Look, therefore, new mountings
were made at the Kapteyn Laboratory.
The first impression of the site on
Klavervlei Farm was sufficiently en-
couraging to make the EC decide on a
more thorough comparison with the
Zeekoegat area. Besides Table Mt.,
some other sites on Klavervlei Farm with
somewhat different local characteristics
were to be investigated: Rockdale Mt.
and Flathill mentioned before. Of the
three sites on Table Mt. only the most
southern one was to be kept. For the
new programme, the Danjon telescopes
were returned to Paris for thorough
overhaul. By the beginning of 1961 they
were available again on the sites.
Meanwhile, in the course of 1960,
plans for Fehrenbach's GPO project had
advanced so far that a search for a
suitable location became desirable. As
we shall see in article IV, it also was at
about this time that the EC agreed in
principle to incorporate the GPO into the
"initial programme" of telescopes men-
tioned in the Convention. For the prepa-
ration of the many logistic measures
connected with it as weil as with the
Klavervlei testing (erection of GPO
housing, satisfactory living quarters for
the observers, water and power supply,
etc), Fehrenbach, Couder and Blaauw
visited South Africa from mid-November
to mid-December 1960. They reported
on their visit at the January 1961 meet-
ing of the EC [18]. As one of the results
of the mission, the choice for the loca-
tion of the GPO fell on the site near
The Comprehensive Programme,
The final, comprehensive programme
was planned to run for at least a full year
but would, in fact, be concluded only in
the course of 1963. It was supervised by
a succession of astronomers, the first
one being again Andre Muller. The
simultaneous monitoring of the four
sites required a larger staff than had
been engaged before, but we realized
that it would by no means be necessary
for all of these to be astronomers. What
we rather needed was: willingness to
spend long periods at isolated spots in
the desert in primitive housing; handi-
ness in technical matters; a gift of im-
provisation and elementary cooking;
and, last but not least, readiness to per-
form over extended periods the routine
work of the testing ...
How to find such people? It occurred
to me that all this sounded like the inter-
ests of an ambitious boy-scout, so we
advertized our wishes in a Dutch journal
of boy-scout leaders. The result was
rewarding: among the applicants were
Albert Bosker and Jan Doornenbal, both
of whom later became employees of
ESO. Among the team that started the
work in March 1961 we also encounter
the two young German astronomers D.
Messerschmidt and W. Schlosser, and
At the Zeekoegat Station, January 1962.
Above: Andre Muller with Bert Bosker, adjusting the mounting of a
Danjon tefeseope, in front of the rondavef that housed the instrumen-
Below: at fett, Ihe "Abri" housing Ihe Marseilles GPO; middle and
right, houses of the observers. Photographs by the author.
At Roekdafe Station on Kfavervfei Farm, January 1962.
Above: B. van Geefen of ZWO, al fett, tafking to observers of the
Tübingen photometrie projeet J. Pfleiderer, U. Haug and Kopp in
Roekdafe farmhouse.
Below: Rockdafe farmhouse on Kfavervfei Farm. Photographs by the
Site testing in the "Comprehensive Program" in 1962. Left: with the Danjon
telescope ror image quality. Right: with photometrie telescopes ror extinclion
measures; here: intercomparison check or two telescopes. Photographs by D.
8eintema, 1962.
G. Bilius, a geodesist from Sweden who
took over the local supervision in May
1961. He was succeeded in this capaci-
ty by H. Linden from Sweden over the
period from August 1961 to April 1962,
and by L. Petterson from Sweden trom
May to October 1962, after which Andre
Muller took over again. Others who over
certain periods collaborated in the site
tests were P. McSharry, a geodesist
trom South Africa, and the young as-
tronomers K. Kopp, W. Seufert, W.
Weber and M. Grewing from the GFR
and D. Beintema from the N e t h e r ~
A working scheme for the operations
had been drawn up by Muller in De-
c e ~ b e r 1960 [19]. From it we quote:
. - - - Irrespective of weather condi-
tlOns, the observers and their assistants
are on duty from sunset till sunrise to do
meteorological and astronomical obser-
vations at regular times.
- - - Observers and assistants have
to work during 25 consecutive nights
and after this period have to take leave
of 5 consecutive nights. These 5 nights,
covering aperiod of nearly 6 days, can
be spent anywhere in the Union of
South Africa and special provisions are
made to meet extra expenses. With the
exception of these 5 nights, there will be
no opportunity for outings, whatsoever.
The groups of observers and assis-
tants will be shifted from one station to
the other at regular times, to ensure a
good comparison between the different
- - - The observers and assistants do
organize their own housekeeping
[whichJ includes foraging in Beaufort
West, Zeekoegat or Prince Albert - - -. "
An interim report on the new tests
was submitted by Muller and Blaauw to
the March 1962 meeting of the EC [20].
They had just returned trom a visit to the
activities in South Africa in December
and January made jointly with B. van
Geelen who, as an associate of J. H.
Bannier of the Dutch organization ZWO,
took care of the many financial, adminis-
trative and personnel matters con-
nected with the site testing. Their report
[21] describes in detail the structure of
the site tests at that epoch. The routine
monitoring of image quality and climatic
conditions proceeded satisfactorily at
the four sites. Of the three on Klavervlei
Farm, Flathill seemed to emerge as the
most favourable one; of the other two,
about equal in quality, it was decided to
lower the priority for Rockdale Mt. The
meeting decided that regular observa-
tions should continue till about March
1963, so that the period on which final
judgement was to be based should con-
tain two complete runs of the normally
most favourable season from November
through March.
Muller reported again at the June and
October 1962 meetings of the EC, after
returning from visits to South Africa. By
October the image quality tests on
Rockdale Mt. had been stopped (as was
the photometric project of Tübingen on
that site, described below). With the ter-
mination of all tests in sight, the EC
appointed a small group to study the
results in preparation for the decision on
the site, to be chaired by Siedentopf and
further consisting of Dommanget,
Fehrenbach, Muller and E. Holmberg
from Sweden, thus having represen-
tatives of the five participating coun-
By the time of this EC meeting of
October 19 and 20, 1962 the ESO Con-
vention had just been signed (October 5)
and the EC took two important
measures. One was the appointment of
Otto Heckmann as acting Director of
ESO, per 1 November 1962 (to be con-
firmed after the ratification). Heckmann
had visited the ESO activities in South
Africa together with Fehrenbach in Au-
gust and September, 1962. Further-
more, by this time the interest of ESO in
the site tests in South America had led
to a mission of Muller and McSharry to
Chile, to join the American group under
J. Stock; McSharry was already on his
way at the time of the October 1962
meeting and Muller was to follow him
shortly thereafter. Their findings will be
reported in the next article.
The routine observations of image
quality and of c1imatic conditions termi-
nated, as had been planned, around
March 1963. At its February 1963 meet-
ing the EC decided to continue the work
in South Africa only for the purpose of
an intercomparison of the Danjon tele-
scope tests with telemeter observations
as they had been used by Stock's group
in the Andes; for this purpose one of
their telemeters was shipped to South
Africa. This final programme was carried
out by McSharry under supervision of
last Tests in South Africa:
the Siedentopf Experiment
Towards the end of the activities in
South Africa, a new kind of test was
introduced that had been developed
over the past years by Siedentopf and
Mayer at Tübingen. It used measure-
ments of the rapid temperature fluctua-
tions which accompany the turbulence
in the atmosphere and which, in turn,
One of the Danjon telescopes, provided in the last stage of the tests
with a mask for experiments by A. B. Muller simulating a double beam
telescope for measuring image motion. Photograph by D. Beintema,
From boy-scout leaders to ESO site testers to ESO employees:
Jan Doornenbal, left, and Bert Bosker. Photograph by D. Beintema,
are correlated to the image quality; see
also the description in the accompany-
ing box. For the measurement of these
temperature fluctuations, thermocou-
pies and resistance thermometers of
small time constant were used. By
mounting the instruments on masts at
different heights, the dependence of
temperature f1uctuation on elevation
above ground level could be measured
and, hence, the dependence of turbu-
lence on height. These experiments
have played an important role in the
decisions taken later with respect to the
level at which the telescopes on La Silla
were to be mounted.
Applications of the method in ESO
context were made by F. Unz in the
period July 17 to September 1, 1963 at
Zeekoegat and Flathill. After verbal pro-
visional accounts by Siedentopf this
work was reported in an ESO publica-
tion in 1964 by (the late) Siedentopf and
Unz [22). Simultaneously with these
measurements, wind velocity was moni-
tored, and the quality of the seeing was
estimated by the measures of image
motion with the double-beam teleme-
ters. Two important results were found:
the amplitude of the temperature fluctu-
ations decreased rapidly with increasing
height within the range of 3.5 to 24
meters, and the difference between the
amplitudes at low and high levels was
strongly correlated with the amplitude of
the image motion. These results im-
mediately led to the conclusion that
mounting the telescopes at high level
above the ground should eliminate most
of the image motion, and hence improve
the seeing. A more extensive report was
published by Unz in 1970 [23).
The Tübingen Photometrie
The Tübingen photometric project on
Rockdale Mt. ran from August 1961 to
November 1962. It was carried out with
a 3-colour photometer on a 40-cm tele-
scope by members of the staff of
Siedentopf: J. Pfleiderer, Miss U. Mayer,
J. Pesch, U. Haug, and J. Dachs. Also,
surface pllotometry was done of the
Milky Way and of the Zodiacal Light in
blue and red. Siedentopf reported at the
October 1962 meeting of the EC on
provisional results. The Rockdale Mt.
observations would become part of the
data later used by the Site-Selection
Committee. Full reports on the Tübingen
project were published in 1966 by
Dachs, Haug and Pfleiderer [24) and by
Pfleiderer, Dachs and Haug [25).
The Marseilles GPO Projeet
Fehrenbach's objective-prism radial
velocity project at Zeekoegat became
fully operational about the middle of
1961, after delays in the construction
phase. It extended in time considerably
beyond the termination of the monitor-
ing of image-quality and climatic condi-
tions, until weil in the year 1966, after
which the GPO was moved to La Silla.
In the first issue of the ESO Bulletin, of
November 1966 (the Bulletins were a
series of ESO publications terminated in
1975), Fehrenbach describes the history
of the GPO work at Zeekoegat. Over the
years, altogether some thirteen col-
laborators of which several with their
families, most of them from the Mar-
seilles staff, had worked on the project.
It produced a large number of GPO
plates, most of them on the Magellanic
Clouds. However, the observational
conditions, although not unsatisfactory,
proved to be inferior to those encoun-
tered soon afterward on La Silla.
One of the disturbing factors were the
strong daily temperature variations at
the Zeekoegat site. Moreover, the GPO
had been mounted at Zeekoegat at
ground level, and one of the measures
taken to obtain better image quality at
La Silla was placing the telescope at
high level. For an early progress report
on the project, see, for instance, the
Information Bulletin of the Southern
Hemisphere, No. 2, Sept. 1962, p. 22.
The Comprehensive Reports on
the South Afriean Tests
We finally arrive at the reports which
sum up the total of the ESO efforts in
South Africa. A final report has been
published as an ESO publication in
1967, long after the decision on the
choice of the site had been taken. It was
compiled at the request of the Director
of ESO by Ursula Mayer of Tübingen
who had participated in some of the
activities in South Africa, and under the
auspices of ESO's Site Evaluation Com-
mittee. It carries the title Astronomical
Site Testing in South Africa and contains
contributions by many of the people
who had participated or had been ac-
tively involved in the tests.
The report systematically surveys
studies of the meteorological condi-
tions, matters of organization, and the
seeing tests and their results in
chronological order. Although the re-
port, in this form, has not played a role in
the decision on the site, it remains an
interesting document, not just for histor-
ical reasons, but also because in its
concise, yet sufficiently detailed presen-
tation it may serve for other purposes of
meteorological and astronomical na-
ture. The booklet is in the ESO Library
(and probably in many institute libraries)
and also forms part of the ESO Histori-
cal Archives [26).
The decisive comparison between
conditions in Chile and South Africa waS
based on provisional reports, but on
virtually the same data as those used for
the final document just mentioned. Tllis
comparison was prepared by Sieden-
topf for the EC Meeting of 15 November
1963; it was published in 1966, in the
first issue of the ESO Bulletin. Sieden-
The Benevolent Environment
In the description of ESO's earliest history we encountered first of all the astronomers
and their immediate collaborators. But their work would not have been possible without
strong administrative support in Europe and the logistic services and hospitality of South
African institutes.
Thraughout the pre-ratification phase, the efforts toward ESO relied heavily on the moral
support from the part of individuals in government departments or in science funding
organizations. For France, the authorities concerned were in the Ministry of Science and
Education and, ultimately, in the Ministry of Foreign Affairs; for the German Federal
Republic and Belgium, in those dealing with science and education or technica/ deve/op-
ment. For the Netherlands and Sweden they belonged to the science supporting organiza-
tion ZWO and the Swedish Natural Science Research Council, respectively. These struc-
tural differences also determined the nature of the sources for the provisional funding of the
site tests.
Particularly meritorious for ESO's early development was the Director of the Dutch
organization ZWO, J. H. Bannier. Fram the moment of his appointment as Treasurer of the
ESO Committee (at its October 1957 meeting) Bannier firmly took in hand the financial
management. His task was not only budgeting and bookkeeping, but also the continuous
effort to persuade the authorities in the partner countries to provide, on the necessarily ad
hoc basis, the required funds. Bannier's authority allowed him, when necessary, to take
initiatives in funding rneasures which made it easier for the partner countries to cross the
financial bridge. From early 1959, Bannier made his associate Dr. B. van Geelen, a young
chemist, available for services inc/uding personnel matters, preparation of travel, insur-
ance, etc. - without frawning upon the bill of a water diviner in South Africa [27].
On the South African side, throughout the site testing there was strong interest and
support fram the part of the President of the Council for Scientific and Industrial Research,
Dr. S. M. Naude. CSIR provided know-how on technical matters required for setting up the
testing stations and made vehicles and measuring instruments available for ESO's rather
demanding use. Responsible for these services was from 1956 CSIR's Director for
International Scientific Relations, Dr. C.G. Hide.
Essential was, of course, the collaboration and support experienced throughout the work
fram the part of the owner of Klavervlei Farm, the Köster fami/y, and of Mrs. Oosthuizen of
Sunnyside Farm at Zeekoegat.
Last but not least, there was the generous hospitality extended to the ESO teams by the
South African observatories. With the testing activities gradually shifting to the Karroo, ESO
relied more and more on the counsel and support provided by the staff of the Cape
Observatory. The outstanding hospitality offered by its Director, R. H. Stoy and Mrs. Stoy,
and by his associate David Evans and Mrs. Evans is warmly remembered by all those who
participated in ESO's South African venture.
topf used the data collected at
Zeekoegat, Flathill and Rockdale Mt.,
and those collected in Chile by Stock
and by Muller and McSharry. We shall
return to it in the next article and men-
tion here only that the report confirmed
what had been strongly suspected: that
the sites in the Andes Mountains around
La Serena were to be preferred on sev-
eral grounds: the number of clear nights,
the image quality, and the surprisingly
lowtemperature drop during the night. It
was at this meeting, 15 November 1963,
that the EC decided in favour of South
In the beginning of this article I re-
ferred to the deterioration of seeing in
the Course of the night in the northern
part of South Africa, mentioned by Bok.
SUC.h systematic change is not explicit-
Iy dlscussed in the reports on the ESO
tests. However, while preparing this arti-
cle, I am informed by Andre Muller that
a!so ,on the Klavervlei and Zeekoegat
sltes this phenomenon was definitely
noted and the deterioration was closely
to the decrease of temperature
In the. Course of the night. In fact,
to Muller, this relation pro-
vldes a strong first indication of the
quality of a potential site: the smaller the
in temperature, the better the site.
Flnally, we note that in the course of
the tests, the rating of image quality by
means of the diffraction rings only was
f.elt. m?re and more as an unfortunate
with "good" rings but
apprecla?le Image motion did OCcur and
were of httle use for practical work Iike
stel.lar spectroscopy, as Was in fact ex-
observers at the GPO.
Thls was POlnted out, for instance in the
report of February 1962 refer;ed to
footnote [21] but it did not lead to
drastlc mOdifications of the techniques
of observation.
At the End, Bewilderment and
rather sudden switch from South
to Chile did not pass without be-
wllderment to the young astronomers
and their collaborators still at work in
Africa. Had years of effort been
walsted? Some disappointment was
Heckmann was aware of
thls and expressed it in a letter to me
which, unfortunately, 1 have not been
able to recover but of which I do re-
member the first words: "Mich drückt
das Bewußtsein ...". Disappointment
Soon make room for the convic-
tlon that the decision had been right.
References and Notes
Abbreviations used:
EHA '" ESO Historical Archives (see the de-
Scription in the previous issue of the
FHA = Files belonging to the Office of the
Head of Administration at ESO Headquarters.
EC = ESO Committee (the committee that
preceded the Council); for a list of the meet-
ings of the EC, see the previous article.
Heckmann Sterne = O. Heckmann, Sterne,
Kosmos, Weltmodelle, Verlag Pi per & Co.,
München, Zürich, 1976.
[1] See the report on the site selection by
W. H. Stevenson and H. Knox-Shaw in
Monthly Notices R.A. S., Vol. 95, p. 447,
[2] In EHA-I.A.1.3. A paper presented at the
F/agstaff Conference on Photo-electric
Problems, Techniques, and Instrumen-
tation, Aug.-Sept. 1952.
[3] In EHA-I.A. 1.3.
[4] H. Siedentopf: Climate of the Union of
South Africa, Astron. Inst. of the Univ. of
Tübingen, 1955, in EHA-1. A. 1.3.
[5] Memo of this meeting in EHA-1. A. 1.3.
[6] A. Danjon and A. Couder, Lunettes et
Telescopes, Paris 1935, Chapitre V. See
also Comptes Rendus No. 183, 1032,
1926 for the calibrations.
[7] EHA-1. A. 1.3. A long report by Elsässer
to Heckmann, Siedentopf and Unsöld
accompanies this letter.
[8] See I. A. 1.5. and I. B. 3.
[9] See I.C.2.3.a.
[10] EHA-I.C.2.3.d.
[11] Comm. Obs. de Belgique (Ucc/e),
No. 141,1958.
[12] See minutes of a discussion on 25 July
1958 following the 8th EC Meeting in
[13] See minutes EC Meeting of July 1958,
item 13 in EHA-I.A. 1.7.
[14] See letter of J. H. Oort to the EC of Oct.
21,1958 in EHA-I.C.2.3.
[15] This report in EHA-1. B. 11. and
[16] See the minutes of this (12th) meeting of
the EC. The report by Muller seems to
be missing from the EHA.
[17] In EHA-I.C.2.5.d.
[18] The report is contained in the minutes of
the meeting.
[19] In EHA-I.C.2.2.a.
[20] In EHA-1. B. 11.
[21] See map EHA-f.A.1.16.
[22] EHA-I.C.2.7.b., H. Siedentopf and F.
Unz, Temperature Fluctuations in the At-
mospheric Ground Layer observed at
Zeekoegat and Flathill (South Africa),
March 1964.
[23] F. Unz, Mitteilungen Tübingen No. 116 =
Meteorol. Rundschau 23, p. 87, 1970.
[24] J. Dachs, U. Haug and J. Pfleiderer, Mitt.
Tübingen No. 87 = J. Atm. Terr. Phys.
28, p. 637, 1966.
[25] J. Pfleiderer, J. Dachs and U. Haug, Mitt.
Tübingen No. 88 = Zeitsehr. für Astroph.
64, p. 116, 1966.
[26] In EHA-I.C.2.7.b.
[27] See letter van Gee/en to B/aauw of 11
November 1960 in EHA-I.C.2.8.d.
Wind Turbulence in the Dome of the 3.6-m Telescope
L. ZAGO and F. RIGAUT, ESO
It is a common experience that tele-
scope tracking may be affected by
strang winds, as in some cases the air
flow and its associated turbulence
penetrate the dome with an amplitude
sufficient to perturbate the smooth op-
eration of the telescope. Reports of this
phenomenon are, however, only qualita-
tive as no measured data were available
(to the authors' knowledge) on quan-
tities such as mean flow penetration,
turbulence intensity and vortex scale in
a dome with an open slit. Only some
recent wind tunnel tests [1] have ad-
dressed this question, although the re-
Iiability of the results may be somewhat
questioned because of the scale similar-
ity problems of wind tunnel simulations
with raund buildings.
Here the results of a preliminary in-
vestigation in the 3.6-m dome are pre-
sented. The purpose of this first series
of measurements was to get a quantita-
tive evaluation of the most critical (worst
case) wind effects in the dome, rather
than a systematic survey representative
of all observing conditions.
The Parameters of Wind
Air flow turbulence may be charac-
terized by several parameters. The most
immediate is the velocity rms Uu mea-
sured along the mean flow direction.
Also one often refers to the turbulence
intensity, which is defined as I = uuiD,
with Dbeing the mean velocity.
Actually the turbulent movement with-
in the mean flow consists of multiple
vortices of different size. A measure of
the mean dimension of these vortices is
given by the so-calied turbulence
scales. The main (longitudinal) scale L ~
is computed from the autocorrelation
coefficient R (x, x, t) of velocity as:
L ~ = DSä R (x, x, 6t) d t
Each size of vortices generates velocity
fluctuations at a given frequency. The
distribution of kinetic energy along the
frequency is given by the spectral power
density function Su(n), computed as the
Fourier transform of the autocorrelation
function. This is often called the gust
(velocity) spectrum and from its defini-
tion it is also:
u ~ = Sä Su(n)dn
The tracking performance of a telescope
is actually affected by the pressure
power spectrum Sp(n) which is obtained
fram a time series of dynamic pressure
values P(t) = ~ gU(t)2, similarly as Su(n)
from U(t). The pressure fluctuations,
represented by Sp (n) (which are seen by
the telescope as forces and moments)
should be compensated by the tracking
control loop. Therefore particularly im-
portant is the amplitude of Sp(n) in the
range beyond the bandwidth of the con-
trol system, typically 1 Hz, which repre-
sents fluctuations which the tracking
loop will often not be able to correct
and, in a worst case, which may even
excite resonance modes in the tele-
Measuring Equipment and
The measurements were taken in the
upper part of the 3.6-m dome, taking
advantage for access of the bridge
crane there iocated, during the evening
of a windy day when the dome
anemometer indicated almost constant-
Iy a mean velocity of 18 m/s from North.
The dome slit was opened as shown in
Figure 2.
A vortex type anemometer (Fig. 1)
was utilized, which is particularly suited
for fast response measurements, having
a 6 t . Uresolution of 6 mm. The mea-
surements consisted of wind velocity
sequences of 137 seconds each with
4096 records, therefore at the frequency
of nearly 30 Hz. Several such se-
quences were recorded at different po-
sitions along the path of the bridge
crane as shown in Figure 2. Before each
sequence, the dome was rotated forth
and back in order to find, rather empiri-
cally, the azimuth angle at which one
would have the stronger feeling of wind
flow and turbulence. Not surprisingly,
this was found to be approximately fac-
ing the mean wind direction. Therefore
the values measured are properly worst
case quantities, as one may expect that
during observation the slit would be fac-
ing the wind only a fraction of the ob-
serving time. The recordings were sub-
sequently processed with the MIDAS
system in order to get statistical para-
Results and Conclusions
The main results from the measure-
ments are given in Figures 3 to 7, in
function of the distance from the edge of
the slit. The data of each figure are
commented with reference to the corre-
sponding parameters of the free wind
flow incident to the dome. Note also that
some data sequences were taken along
the centre of the slit, others along the
left side: this simply because we re-
marked that flow and turbulence were
Figure 1: The vortex anemometer placed near the slit (a) and further inside the dome (b).
Figure 2: Measurement loeations in the 3.B-m dome.
Open section of slit
during the measurements
t:. Center of slit
+ Left side
:j: (f)
o 2 4 6 8 10
3: Mean flow velocity D. The mean
wmd velocity outside the dome was about
18 m/s. Already just inside the slit this is
redueed to about 5 m/s. However fur-
ther inside the dome centre, one still reeords
up to 3 m/s.
often stronger along the sides than
along the slit midline.
We have verified that the 3.6-m dome,
even with the slit facing the wind, acts
as an efficient wind shield in terms of
mean flow velocity. Nevertheless, the
slit is the cause of velocity fluctuations
inside the dome, which are definitely
larger than in the original atmospheric
turbulence. This dome induced turbu-
lence has a mean scale of the order of
the slit width and a peak frequency in
the range 0.3 to 1 Hz. In proximity of the
slit this effect causes also pressure vari-
ations which are larger than in the free
cjc 6
o 2 4 6 8 10
DISTRNCE FROM SLiT "''1)
Figure 5: Turbulenee seale L,J. The turbulenee
seale in the free atmosphere at the level of
the 3.B-m dome is likely to be in the range 50
to 200 metres. This free flow turbulenee is still
a eontributing faetor near the edge of the
dome (L,J = 14-17 m); further inside we have
purely slit made turbulenee, with an average
seale of the order of slit width.
atmosphere, particularly in the frequen-
cy range above 1 Hz where they might
directly affect the tracking behaviour of
an hypothetic telescope whose top
structure would come closer to the slit
than the present 3.6-m one. Further in-
side the dome, because of the large
decrease of mean velocity (note that,
from the definition of dynamic pressure,
op oc D0u + the amplitude of
pressure oscillations is largely below the
situation in the free flow.
When dealing with flow turbulence
around or inside telescope domes, a
question which is often raised is
whether this is Iinked to the thermal
microturbulence causing dome seeing.
Although the measurements described
here meant to address only the problem
of wind disturbance on tracking, the evi-
dence found of important and large flow
Free atmospheric f10w
-------------- 1 m from edge, slit center
-"---'-'-'-'-'" 5 m from edge, slit center
-------- 9 m from edge, slit center
------ 3 m from edge, left side
- - - - - - - - - 9 m from edge, left side
Figure 6: Normalized gust speetra at different loeations inside the dome and, for referenee, in
the free flow. In the free atmosphere most of the wind turbulenee energy is found in the range
0.01 to 0.05 Hz. In the dome a peak in the range 0.3 to 1 Hz appears, whieh takes more of the
turbulent energy the further away one is from the slit. Note that the speetra are here normalized
with the respeetive values.
+ 6
o 2 lj 6 8 10
OlSTANCE FROM SLlT [MJ
Figure 7: Turbulent pressure high-pass fil-
tered at 1 Hz:
(Jp>lHz = Vf/Hz Sp(n)dn
This quantity is an approximate indicator of
the dynamic wind loading which, if acting on
an hypothetic telescape, cannot be compen-
sated by the tracking control loop, assumed
here to have a bandwidth of 1 Hz. The refer-
ence free flow value is 3.2 Pa: one may note
that inside the dome this value is exceeded
only in proximity of the slft. Further inside, the
amplitude of pressure fluctuations is largely
below the level in the free atmosphere.
vortices quite deep inside the dome may
have relevance also to dome seeing, in
particular to the energy balance of the
Visiting Astronomers
(April1-0ctober 1, 1989)
Observing time has now been allocated for
Period 43 (April 1-0ctober 1, 1989). The
demand for telescope time was again much
greater than the time actually available.
The following list gives the names of the
visiting astronomers, by telescope and in
chronological order. The complete list, with
dates, equipment and programme titles, is
available from ESO-Garching.
3.6-m Telescope
April: Oosterloo/van der Kruit, Danziger/
CappellarolTuratto, di Serego Alighieri/
Tadhunter/Fosbury, RuiziMaza, Renzini/
D'Odorico/Greggio/Bragaglia, Danziger/
Moorwood/Oliva, Moorwood/Oliva, Danzi-
bury/Fransson, Krautter/Starrfield/Ögelman,
May: BertolaiBuson/Zeilinger, Krautter/
Starrfield/Ögelman, ScaramellaiChincarinii
VettolanilZamorani, Ilovaisky/Chevalier/
Pedersen, Surdej et al. (2-003-43 K), Schmi-
der/FossatiGrec/Gelly, Butcher, Molaro/
Spite F.Nladilo, Butcher/Pottasch/Slinger-
land/Baade/Christensen-D.lFrandsen, Dan-
bury/Fransson, Seggewiss/MoffatiRobert.
June: Seggewiss/MoffatiRobert, Cacciari/
Clementini/PrevotiLindgren, Piotto/Ca-
paccioli, Pasquini, CacciarilClementini/Pre-
votiLindgren, Perrier/Mariotti/Mayor/Du-
quennoy, Pottasch/Pecker/KarojilSahu K. C.,
Particularly in large domes, it is una-
voidable that temperature gradients ex-
ist between different sections and struc-
tures. Then any flow turbulence will in-
crease the heat transfer between dome
and air, therefore feeding energy to the
microthermic turbulence which causes
dome seeing. In this respect the dome
mechanical turbulence may act as an
intensifier of the known seeing effect
created by temperature differences be-
tween air and surfaces or between in-
side and outside conditions.
The authors wish to thank J.-L
Sauvageot and C. Santini of ESO La
Silla for helping to carry out the mea-
surements. They are also grateful to the
dome maintenance staff for their coop-
eration and apologize for some incon-
veniences involuntarily caused during
the tests.
[1] Kiceniuk T. and Potter K., Internal Air Flow
Patterns for the Keck 1O-m Telescope Ob-
servatory Dome, GALCIT Report 10'
July: Häfner/Barwig/Schoembs, Augus-
teijnlvan Paradijs/Schwarz, Meylan/Mayor/
Dubath, Weidemann/Koester/Jordan, Ar-
deberg/Lindgren H.lLundström, v. d. Veen/
Habing/Geballe, Krabbe/Sternberg, Veron/
Hawkins, de Lapparent et al. (1-003-43 K).
August: de Lapparent et al. (1-003-43 K),
Dubath, Bergvall/Rönnback, Guzzo/Collins/
Heydon-Dumbleton, Neri/Grewing/Bässgen
M., Danziger/BouchetiGouiffes/LucylWamp-
ler/Fosbury/Fransson, Wagner, Leibundgutl
Tammann, lovino/Shaver/CristianilClowes/
September: lovino/Shaver/Cristiani/
Clowes/Pierre, Rhee et al. (1-005-43K), Ma-
zure et al. (1-014-431<), Ögelman/Gouif-
fes/Hasinger/Pietsch/Pedersen, Danziger/
Fransson, Webb/Carswell/Shaver, Bergeron/
D'Odorico/Petitjean, Marano/CappilHeld,
Schwope/Beuermann, Miley et al. (2-001-
43 K), Bender et al. (1-004-43 K), Bergeron et
al. (1-012-43 K).
2.2-m Telescope
April: Oosterloo/van der Kruit, Gallettal
Bettoni, Katgertlvan Haarlem, Giraud,
Srinivasan/Danziger, Reinsch/Beuermann/
Festou/Pakull, Meisenheimer/Röser, Moneti/
Zinnecker/Reipurth, Moorwood/Oliva, Dan-
bury/Fransson, Tadhunter/Pollacco,
Reinsch/Beuermann/Festou/Pakull, Cetty-
May: Cetty-VeronlWoltjer, Bertola et al.
(1-008-43 K), Surdej et al. (2-003-43 K),
Srinivasan/Danziger, Bertola et al. (1-008-
June: MPI TIME, ChinilWargau, Glass/
Moorwood/Moneti, Ortolani/Piotto, Piotto/
July: Brahic/Sicardy/Roques/Barucci,
Habing/Le Poole/Schwarz/van der Veen,
Brahic/Sicardy/Roques/Barucci, v. d. Veen/
Habing/Blommaert, v.d. Veen/Habing/
Geballe, Tosi/Focardi/Greggio.
August: Richtler/Kaluzny, Wiklind/Bergvall/
Aalto, Bergvall/Rönnback, Tanzi/Bersanelli/
BouchetiMaraschi/FalomolTreves, Danziger/
Fransson, Jörsäter/Bergvall, Appenzeller/
Wagner, Mi/ey et al. (2-001-431<), Christen-
September: Barbieri et al. (2-007-43 K),
Bender et al. (1-004-43 K), MPI TIME.
1.5-mSpectrographic Telescope
April: Danziger/BouchetiGouiffes/Lucy/
Wampler/Fosbury/Fransson, ChincarinilDe
Souzaldi Stefano/Sperandio/Molinari, Cour-
voisier/Bouchet, Danziger/BouchetiGouiffes/
LucylWampler/Fosbury/Fransson, Buzzo-
Fosbury/Fransson, ThelWesterfundNardyal
de Winter, Danziger/BouchetiGouiffes/Lucy/
Wampler/Fosbury/Fransson, Andreae/Drech-
May: Tadhunter/Pollacco, Danziger/
Fransson, LanziArtru, Gehren/Steenbock!
Rei/e/Axer/BurkertiFuhrmann, Spite F.ISpite
M., Danziger/BouchetiGouiffes/LucylWamp-
ler/Fosbury/Fransson, Acker/Stenholm/
June: Acker/Stenholm/Lundström, Baade/
Stahl, Gerbaldi et al. (5-004-431<), Waelkens/
LamerslTranslWaters, Pottasch/Pecker/Ka-
July: Courvoisier/Bouchet, BicaiAlloin, v.
Genderen/v.d. HuchtiSchwarz/de Loore,
Baribaud/Alioin/PelatiPhillips, Boffin/Joris-
sen/Arnould, Hron, Baribaud/Alloin/Pelatl
Phi/li ps, Wiklind/Bergvall/Aalto.
August: Danziger/BouchetiGouiffes/Lucy/
Wampler/Fosbury/Fransson, Baribaud/AI-
loin/PelatiPhillips, Tanzi/Bersanelli/Bouchetl
Maraschi/FalomolTreves, Danziger/Bouchetl
Eriksson/Gustafsson/Olofsson, Danziger/
Fransson, Baribaud/Alioin/PelatiPhillips,
KatgertlRhee, Baribaud/Alioin/PelatiPhillips.
September: Danziger/BouchetiGouiffes/
LucylWampler/Fosbury/Fransson, Cappil
ChincariniNettolani, Baribaud/Alloin/Pelatl
Phillips, Danziger/BouchetiGouiffes/Lucy/
Wampler/Fosbury/Fransson, JugakulTeka-
da-Hidai/Holweger, Gerbaldi et al. (5-004-
431<), Baribaud/Alloin/PelatiPhi/lips, Danzi-
bury/Fransson, CalvanVMarziani, Baribaud/
April: Franco, Baade/v. Kerkwijk!Waters/
Henrichs/van Paradijs, Gratton/Gustafsson/
Eriksson, Westerlund/Krelowski, Mathys,
May: GilletiCrowe, Baade/v. KerkwijkIWa-
ters/Henrichs/van Paradijs, Stalio/Franchinil
Spite F./Spite M., Crane/Palazzi/Mandolesil
June: Crane/Palazzi/Blades/Kutner,
HUbert, Danks/Crane/Massa, Pasquini,
Gredel/v. DishoeckiBlack.
July: Pottasch/Sahu, Diesch/Bässgen M./
Grewing, Didelon.
August: de Vries/van DlshoeckiHabing,
September: Foing/CrivellariNiadilo/Cas-
telli/Beckman/Char/Jankov, Prein/van Gen-
deren/Zwaan, Gustafsson/Eriksson/Olofs-
son/LambertiParesce, Thimm/Hanuschiki
1-m Photometrie Teleseope
April: Lorenzetti/BerrillilCeccarelli/Nisinil
Saraceno, Lorenzetti/Ceccarelli/Liseau/Ni-
sini/Saraceno, Scaltriti/BussolOrigliaiDe
Silvestro, Schultz, Courvoisier/Bouchet Dan-
bury/Fransson, Catalano F.A./Kroll, Madejs-
Madejsky/Appl, Gouiffes/Cristiani,
Remsch/Pakuli/Festou/Beuermann Cour-
vOisier/Bouchet, Danziger/BoucheVGouiffesl
LucYlWampler/Fosbury/Fransson, Kreysingl
Kaelble/Grewing, v.d. HuchtIThelWilliams
PrevotiLindgren H. '
June: Hahn/LagerkvistiMagnusson/Lind-
gren M., Cacciari/Clementini/PrevotiLind-
gren H., Terzan, WinkiGreve, Courvoisierl
Bouchet, ManfroidNreuxiGosset.
July:. Brahic/Sicardy/Roques/Barucci,
GOUiffes/Cristiani, Schneiderl
August: Poulain/DavoustiNieto, Erikssonl
DI Martmo/ZappalaiCellino/FarinellaiDavis
Alcaino/LilierlAlvaradolWenderoth. '
September: Alcaino/LilierlAlvaradolWen-
deroth, Barbieri et al. (2-007-43K), Danzigerl
Fransson, ZickgraflWolf, Gouiffes/Cristiani.
50-em Photometrie Teleseope
April: Scaltriti/BussolOrigliaiDe Francescol
Schultz, PorettilAntonello, ThelWesterlundl
de Winter.
May: ThelWesterlund/de Winter, Staliol
Covino/Neri, Seggewiss/MoffatiRobert.
June: Seggewiss/MoffatiRobert, Kohou-
tekIWenskat, Houdebine/Panagi/Foing/But-
ler/Rodono, Carrasco/Loyola.
July: Group for Long Term Photometry of
August: Sinachopoulos, Foing/Crivellaril
Carrasco/Loyola, Foing/CrivellariNiadilol
September: Foing/CrivellariNladilo/Cas-
telli/Beckman/Char/Jankov, Group for Long
Term Photometry of Variables.
GPO 40-em Astrograph
April: Eist, Scardia.
May: Scardia, Landgraf.
June: Landgraf, Aniol/Seitter/Duerbecki
July: Aniol/Seitter/DuerbeckITsvetkov.
August: Debehogne/Machado/Mouraol
CaldeiraNieiraiNetto/ZappalaiDe Sanctisl
September: Debehogne/Machado/Mou-
rao/CaldeiraNieiraiNetto/ZappalaiDe Sanc-
tisl LagerkvistiProtitch-B./JavanshirlWosz-
1.5-m Danish Teleseope
April: DANISH TIME, Olsen, Della Vallel
Ortolani/Fusi Pecci/Buonanno/Renzini/Fer-
May: OrtolanilFusi Pecci/Buonanno/Renzi-
ni/Ferraro, de Jong/Slijkhuis/Hulvan der
Bliek, Gregorini/MessinaNettolani, 1I0vaiskyl
Chevalier/Pedersen, DANISH TIME.
June: DANISH TIME, CalvanilD'Odoricol
Zwitter, Reinsch/Pakuli/Festou/Beuermann,
Calvani/D'Odorico/Zwitter, Mayor et al. (5-
001-43 K), DuerbeckIVogtiLeibowitz, Bandie-
ralvan den Bergh.
July: Bandieralvan den Bergh, Reinschi
Pakuli/Festou/Beuermann, AzzopardilLe-
queuxlRebeirot, Gratton/Ortolani, DANISH
August: DANISH TIME, Mayor et al. (5-001-
43 K), Ardeberg/Lindgren H./Lundström,
Meylan/Mayor, Bender et al. (1-004-43 K).
September: Bender et al. (1-004-43 K), de
Goudfrooij, Vettolani/CappilGarilli/Gregorinii
Maccagni, Ardeberg/Lindgren H./Lundström,
50-em Danish Teleseope
April: DANISH TIME.
May: DANISH TIME.
June: Ardeberg/Lindgren H./Lundström,
July: DANISH TIME.
August: Group for Long Term Photometry
of Variables.
September: Ardeberg/Lindgren H./Lund-
gO-em Duteh Teleseope
April: DUTCH TIME, van Genderen.
May: van Genderen/v.d. Huchtlvan
Genderen, DUTCH TIME.
June: v. Amerongen/v. Paradijs.
July: v. Amerongen/v. Paradijs, van Para-
dijs/Strom/van der Klis/Spijkstra, v. Gende-
renIv. d. HuchtiSchwarz/de Loore, DUTCH
August: DUTCH TIME, Preinlvan
September: van Genderen, v.d. Huchtlvan
Genderen, DUTCH TIME.
61-em Boehum Teleseope
April: Lemmer/Dachs, Schneider/Jenknerl
May: Schneider/Jenkner/Maitzen.
May: SWEDISH TIME, Israel, de Graauw,
Danziger, Reipurth, van der Veen, Schwarz.
June: SWEDISH TIME.
July: Henkel, Wielebinski, Israel, Eckart,
Becker, Zinnecker, Israel, Henkel, Zinnecker,
Omont, Pottasch, Chini.
August: SWEDISH TIME.
September: Harnett, Combes, Dupraz,
Dennefeld, Bujarrabel, te Lintel, Hekkert, Le
Bourlot, Haikala, Brand, Wilson, Roland.
IRe + 10216: a Peanut Nebula!
T. LEBERTRE, P. MAGAINand M. REMY, ESO
1. Carbon Stars with Shells
Carbon stars with low effective tem-
perature (2,000-3,000 K) are thought to
be long-period variables evolving on the
Asymptotic Giant Branch (AGB). These
objects are burning alternately hydrogen
and helium in different shells around a
degenerate core of carbon and oxygen
[1]. Material processed during the
helium burning phase is dredged-up by
convection to the surface and enriches
it in carbon relative to oxygen.
Objects on the AGB are loosing mass
due to a combination of two processes:
pulsation of the central star and radia-
tion pressure on grains. Consequently,
carbon stars are surrounded by shells
which also have a carbon-rich com-
position. The dust which is formed in
these shells is expected to be mainly
carbon-rich. Its composition is still a
matter of controversy: graphite or
amorphous carbon are generally pro-
posed, but also silicon carbide (SiC) and
magnesium sulfide (MgS). Depending
on the optical depth of the circumstellar
dust shell (CDS), the central star may or
may not be observable. In the first case,
one speaks of a carbon mira and, in the
second case, of an extreme carbon star
(ECS). In fact, it can be shown that there
is a continuity between carbon miras
and ECS's, and that the latter are only
extreme miras undergoing mass 1055 at
a huge rate [2].
One of the most interesting carbon
stars is IRC + 10216. This object was
discovered as an infrared source in a
sky survey at 2.2 11m; it owes its name
(IRC) to this circumstance. Early studies
[3] showed that it is variable with a
period of - 600 days and that it appears
extremely red with a colour temperature
of 650 K over the range 1to 20 11m; such
an energy distribution can only be
understood if the central star is
surrounded by an optically thick CDS
which absorbs stellar radiation and re-
radiates it at longer wavelengths. Its dis-
tance from the Sun is evaluated to be
- 200 pc. Many molecules have been
detected at radio wavelengths in its cir-
cumstellar shell and, from modelling of
the CO emission, a mass 1055 rate of
- 10-
has been derived. Being
so near to the Sun and undergoing mass
1055 at such a large rate, this carbon-
rich source is one of the best studied
and is often considered as the prototype
of ECS.
2. The Shell of IRe + 10216
Its proximity has allowed spatial reso-
lution of its shell. Already in 1969, its
optical counterpart was noted to be
diffuse and elongated at position angle
(PA) - 30°[3]. This diffuseness is main-
Iy produced by scattering of stellar
photons in the CDS; however, especially
at short wavelengths (I.. < 0.5 11m),
photons scattered from the interstellar
radiation field mayaiso contribute. At
infrared wavelengths (2-10 11m), the
source is observed to be extended, with
a typical size of the order of 1" or less; in
this spectral range, thermal emission by
dust in the CDS is dominating. Using
one-dimension speckle-interferometric
techniques, the source is seen elon-
gated in the North-South direction; fur-
thermore, at 211m, it appears asymmet-
rical, being more extended towards the
North and the North-East [4]. A large
polarization is also observed in the opti-
cal as weil as in the infrared ranges. This
information is generally interpreted in
terms of an axisymmetrical structure
with an equatorial disk and polar lobes
like in bipolar nebulae.
Although the bipolar nebula hy-
pothesis is attractive, doubts have been
cast on it. Molecular-line observations
do not present any evidence of devia-
tion from spherical symmetry on scafes
between 10" and 60". Also, deep im-
ages of IRC + 10216 were obtained
through filters Gunn g, r, i and z, in April
1987, with the 2.2-m telescope at La
Silla. They showed an extended struc-
ture up to at least 10" with no evidence
ofaxisymmetry [2]. Furthermore, the
central part « 1") was seen elongated
like in earlier images, but at a PA of
340° instead of 30°. Such a change of
PA in less than 20 years is not easily
reconcilable with abipolar geometry
which presumes a stable structure like
an equatorial disko Finally, the broad-
band energy distribution of IRC + 10216
between 0.5 11m and 3 mm is weil
understood in terms of a radiative trans-
fer model, consisting of a central star
surrounded by a spherical CDS [5]; this
result tends to support the idea that the
geometry cannot deviate too much from
To reconcile the evidences of spheric-
ity (at least on large scales) with the
evidences of asymmetry given by imag-
ery of the central part (on scales < 1"), it
has been proposed that mass 1055
occurs in an inhomogeneous manner
with no systematic trend [2,4]. Convec-
tive cells at the surface of the star may
be very large and induce mass 1055, at a
given instant, in a preferential direction.
Due to stellar rotation, they move with
respect to the CDS; also, with time
these cells evolve, and the matter in the
shell appears clumpy with no systema-
tic deviation from spherical symmetry.
In Figure 1, a schematic representa-
tion of the inner part of the IRC + 10216
CDS, based on such a proposition, is
presented. Matter has been flowing
away preferentially from an active spot
during the last· 20 years. Dust is con-
densing out of the gas when it reaches a
distance, Re, such that grains can sur-
vive in the radiative field; this distance
represents the internal radius of the
CDS and defines a spherical cavity in-
side which no dust is present. If conden-
sation temperature is around 1,000 K,
the apparent radius of this central cavity
is - 0':15, whereas the apparent radius
of the star is - 0':02 (for a stellar effec-
tive temperature of 2,200 K). The dust
density is enhanced in a preferential di-
rection and the images at all wave-
lengths (i.e. in light scattered or emitted
by dust, independently of optical depth)
are observed to be efongated in that
Also, in this scheme, the direction of
polarization is always perpendicular to
that preferential direction as is observed
on IRC + 10216. Due to stellar rotation,
the elongation, which was observed to
be at PA -30°in 1969, appears now at
PA - 340°; dust which reached con-
densation point twenty years aga has
travelled a distance, R
- 3.0 10-
and is now at - 0':5 from the central
star. Finally, if this representation is
correct, the star should appear offset to
the South of the nebula; this conse-
quence could give a natural explanation
to the asymmetry of the profiles ob-
tained at 211m.
3. High-Resolution Imaging
Such a scheme offers an interesting
and plausible alternative to the bipolar
nebula hypothesis, but is still highly
speculative; imagery at high spatial res-
olution (-0':1 or better) would be
necessary to confirm it. The images ob-
tained at the 2.2-m in April 1987 were
acquired in good seeing conditions
(FWHM of stellar images< 1") but their
spatial resolution was largely insuffi-
cient. However, two elements allow to
consider improving their quality by num-
erical treatment. First, a star is close to
IRC + 10216 (-36" to the South-East)
and could be registered simultaneously
on the same CCD frames (see Figs. 1
and 2 in [2]). Second, the pixel size was
small enough (0:"6) so that the images
of this point source were weil sampled.
Therefore, it appeared possible to im-
prove the resolution of these images
with a deconvolution method using as
point spread functions the profiles of the
nearby star. In Figure 2, the images de-
convoluted by the method of maximum
entropy [6] are presented.
The result is striking: one sees a
peanut-shaped nebula whose main axis
is oriented North-South and concavity is
turned towards East. The resemblance
Figure 1: Schematic representation ot IRC +
10216 dust shell (adapted trom [2J). R
is the
radius of the central star (-0:'02). Re is the
inner radius ot the circumstellar dust shel/
(-0:'15); it corresponds to the distance at
which grains are condensing out of the cir-
cumstellar gas. R
indicates the current posi-
tion of matter which reached the condensa-
tion point twenty years ago. North is up and
East at right.
Figure 2: IRC + 10216 images obtained in April 1987 through Gunn g, r, z and i (clock-wise
startmg {rom the upper left cornerl and deconvoluted as explained in the text; pixel size is 0:'26.
Th '. I
e ongmal Gunn rand z {rames were presented in [2j. North is up and East at right.
the COS. The observational confirma-
tion of the scheme proposed earlier by
us [2] is important not only because it
supports our theses, but also because it
contradicts the bipolar hypothesis and,
therefore, the models relying on it. It
definitively pi aces IRC + 10216 among
the normal carbon-rich miras, and not,
as sometimes suggested, among the
protoplanetary nebulae.
From a more general point of view, it
can be noticed that a lot of information
is often present in astronomical data
which is not exploited. An illustration of
overlooked information was given re-
cently in the Messenger [8]; our work
gives another example. It would surely
be worth to apply the same kind of
technique to the IRC + 10216 images
obtained in the optical range during the
last 20 years. Such an investigation
could give indications on the recent
mass loss history and on the stellar rota-
tion period. Moreover, our work demon-
strates the interest in good seeing and
good sampling of the point spread func-
4. What is in a Name?
with the sketch presented in Figure 1 is
so flagrant, that we ourselves could not
believe it. As IRC + 10216 is more easily
observed around maximum we decided
t . '
o walt for the next one (expected to
OCCur around November 1988 [2]) and to
perform new observations at such
Phase. (with a smaller pixel size) in order
to venfy the meaningfulness of this re-
sult. In the meantime, the algorithm of
deconvolution and the d
checked' " . co e were
. uSlng artlflclal images Also
vanous test .,
s on convergence consis-
tency et· ,
. ,c. uSlng the real IRC + 10216
Images were performed; all resulted
positive. For example, one of these tests
conSlsted in rotating an original image
by an arbitrary angle, deconvoluting it
then t t' '
ro a Ing the deconvoluted image
backwards and comparing it to the im-
age obtalned by direct deconvolution; in
Figure 3: IRC + 10216 image obtained in Oe-
cember 1988 through Gunn i" pixel size is
0"17' '
. 5. North IS up and East at right.
all cases, the comparison was satisfac-
tory. Finally, another image-restoration
method [7] was used and gave similar
IRC + 10216 was reobserved on Oe-
cember 16/17, 1988 using the 2.2-m
equipped with the recently com-
misioned CCO no. 15. The individual de-
tector size is 15 !",m which translates to
a pixel size of 0':175 on the sky. The
infrared monitoring performed at the 1-
m was indicating that we had, as pre-
dicted, just passed the maximum by a
few days and that the source was
undergoing a maximum brighter than in
April 1987. We were therefore expecting
slightly larger images for IRC + 10216.
On the other hand, the image quality
was not as good, being around F3, as
measured on the nearby star. The
frames obtained through the same Gunn
g, r, i and z filters were reduced using
the standard procedures; the Gunn i
image is presented in Figure 3. No basic
difference can be noted with respect to
the images obtained in April 1987 [2].
The deconvoluted Gunn i image is
presented in Figure 4. The same struc-
ture as in Figure 2 is again clearly seen.
The similarity between the April 1987
deconvoluted images and the Oe-
cember 1988 ones leads to the convic-
ti on that the peanut shape is not an
observational artifact. Moreover, the
fact that this structure is observed at
600 days difference indicates also that it
is not a transient, but (on a time-scale of
one stellar cycle) a permanent feature of
Finally, astronomers are used to give
fancy names to their pet objects (de-
rived from their food habits?), e. g.: the
Egg Nebula, the Hamburger Nebula
and, even, the Rotten Egg Nebula! For
IRC + 10216, we propose: the "Peanut
Nebula"; as its circumstellar shell is
known from observations in the radio
range to be rich in organic molecules,
we hope this denomination will also
satisfy our colleagues from radio as-
Figure 4: Same as in Figure 3, but deconvo-
luted as explained in the text.
Guidelines for Authors of Articles for the ESO Messenger
[1] Iben, 1., Jr.: 1981, in Physical Processes
in Red Giants, I. Iben Jr. and A Renzini
(eds.), D. Reidel Publishing Company 3.
[2] Le Bertre, T.: 1988, Astron. Astrophys.
(3] Becklin, E. E., Frogel, J.A, Hyland, AR.,
Kristian, J., Neugebauer, G.: 1969, Astro-
phys. J. 158, L 133.
(4] Dyck, H. M., Zuckerman, B., Howell, R R,
Beckwith, S.: Publ. Astron. Soc. Pac. 99,
[5] Le Bertre, T.: 1987, Astron. Astrophys.
176, 107.
(6] Skilling, J., Bryan, R. K.: 1984, Month.
Not. Roy. Astr. Soc. 211, 111.
(7] Magain, P.: 1989, in preparation.
[8] Heydari-Malayeri, M.: 1988, The
Messenger 53, 23.
List of ESO Preprints
December 1988 - February 1989
625. (1) C. N Tadhunter, RAE. Fosbury, S.
di Serego Alighieri: Beamed lonizing
Radiation in Radio Galaxies.
(2) R Morganti et al.: What Are the
Emission Line Filaments Along the
Radio Axis of Centaurus A?
Papers presented at the Como Work-
shop on BL Lac Objects: 10 Years After
(September 1988).
626. G. Garay, R Gathier, L. F. Rodriguez:
Radio Recombination Line Observa-
tions of Compact Planetary Nebulae.
Astronomy and Astrophysics.
627. I. F. Alania, O. P. Abuladze, R M. West:
uvbyß Photometry of Peculiar Band A
Stars, Discovered at Abastumani. As-
tronomy and Astrophysics, Supplement
628. E. Gosset et al.: Analysis of the Photo-
metrie Variability of WR 40. Monthly
Notices of the Royal Astronomical Soci-
629. P. Molaro et al.: Nallnterstellar Absorp-
tion in the Direction of Two LMC Super-
giants in the Field of SN 1987A As-
trophysical Journal, Letters.
630. G. Meylan, S. Djorgovski: UM 425: A
New Gravitational Lens Candidate. As-
trophysical Journal, Letters.
631. P. Crane et al.: Cosmic Background
Radiation Temperature at 2.64 mm. As-
trophysical Journal.
632. R. M. West, H. E. J0rgensen: Post-
perihelion Observations of Comet
Halley at r = 8.5 AU. Astronomy and
633. (1) A. F. M. Moorwood: Near IR Spec-
troscopy of Active Galaxies.
(2) A.F.M. Moorwood, E. Oliva: In-
frared (Fell], Hand H
Lines in Galaxy
To appear in Proceedings of the 22nd
ESLAB Symposium, Infrared Spectros-
copy in Astronomy. ESA SP-290
(AC.H. Glasse, M.F. Kessler and R.
Gonzalez-Riesta eds.).
634. E. Covino et al.: EE Aquarii: a Marginal
Contact System. Monthly Notices of
the Royal Astronomical Society.
635. G. Contopoulos et al.: Comparison of
Stellar and Gas Dynamics of a Barred
Spiral Galaxy. Astrophysical Journal.
The Messenger is ESO's house-journal
and serves as a link between ESO and the
user community. It brings information about
scientific and technical developments at
ESO and also about administrative
measures. At the same time it aims at pro-
viding interesting news about astronomy
and astrophysics to a broader public, in-
eluding policy makers, science teachers,
amateur astronomers, inside and outside
the ESO member countries. At the present
time, the Messenger is distributed free of
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overseas destinations).
The Messenger is abstracted by several
services, including AAA.
Since there are limits to the size of each
issue, set by the available manpower and
by the budget, it has now become
necessary to establish a minimum of
guidelines for authors. They are not in-
tended to restriet the information flow; on
the contrary, they aim at keeping the
"fresh" look of the Messenger, by ensuring
that each issue carries a broad variety of
"interesting" and informative articles.
The Messenger normally appears at the
beginning of March, June, September and
December. The corresponding deadlines
are 6 weeks before, i. e. on January 20, April
20, July 20 and October 20.
Last-minute, "spectacular" (and short)
news items can be received during the
editorial process; please contact the editor
Contributions must be written in English.
For the benefit of our Spanish-speaking
readers, condensed versions of some arti-
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Normally, the editor will solicit contribu-
tions to the ESO Messenger by writing to
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next deadline. Unsolicited manuscripts are
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Submitted manuscripts will be checked
for obvious, technical errors as far as possi-
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revisions of the language. Therefore, if you
are worried about your English, please ask
a colleague to help, before you submit your
Messenger articles are normally not ref-
ereed; however, in certain cases the editor
may solicit the advice of other ESO as-
tronomers before accepting an artic/e.
On rare occasions, especially if there are
important news items which necessitate
last-minute layout revisions, an article will
1. R. N. Wilson et al.: Active Optics 111:
Final Results with the 1m Test Mirror
and NTT 3.58 m Primary in the Work-
have to be delayed to the next issue. The
authors will be informed about this immedi-
ately and will have the opportunity to revise
the artiele, if so desired.
Text and Style
All articles brought in the Messenger
must have some connection to ESO. They
will often be based on results from observa-
tions at La Silla, but mayaiso concern
developments elsewhere of direct implica-
tion for ESO programmes.
An article should normally not exceed 4
printed pages, ineluding figures, but shorter
contributions, down to a picture with an
appropriate caption, are of course most
welcome. One printed text page is roughly
equivalent to 3.5 double-spaced A4 type-
written pages, 7,000 characters or 1,200
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therefore about 10 A4 typewritten pages
(20,000 characters or 3,500 words), plus
The style should be light, but informative.
The Messenger is no substitute for the
professional journals and its articles
should contain less detail and more
background than what is usual in scientific
papers. Remember that the Messenger is
read by many people from different fields
and with a range of background knowledge
from the amateur to the specialist.
A "personal touch" in the form of an
associated event, an unusual result, etc.,
will be appreciated. Similarly, a Messenger
artiele may contain information which is not
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instance about technical problems, rather
speculative ideas, suggestions for future
work, advice to other observers, etc. Early,
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come, but cannot serve to replace proper
accounts in referee journals.
An artiele should not contain more than
10-15 references. Please use the style of
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A maximum of 6 figures is normally
allowed per artiele. They should be submit-
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Colour pictures may be used, if the sub-
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No reprints are made of Messenger arti-
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request should be made when the article is
shop. Journal of Modern Optics.
2. B. Delabre et al.: Astronomical Spectro-
graph Design with Collimator Compen-
sation of Camera Chromatism (4C).
Proceedings of SPIE No. 1055.
A Distant View of La Silla
W. C. KEEL, University of Alabama, USA
The domes and related structures on
La Silla are prominent features of the
Atacama landscape as seen for many
kilometres, and would be for many more
were it not for the intervening Andean
foothills. Under proper conditions, they
may be seen (with the unaided eye) for
even greater distances. Possibly the
most distant normally occupied vantage
point (barring a SPOT image) is from the
Inter-American Observatory on Gerro
Tololo, about 103 km south-southwest
of La Silla. For month-Iong periods twice
a year (when the sun's declination is
about -10°) the domes at ESO are
prominent shortly before sunset, as
shown in tlle photograph taken in mid-
October of 1988. Five structures are
prominent with another two visible on
the original print. Part of La Silla is
hidden behind a foreground mountain,
identified in t11e Mapa Fisico de Chile as
Gerro EI Pozo (east of Almirante Lator-
re); its presence accounts for the greater
difficulty of locating Gerro Tololo from
La Silla, even using binoculars.
It is of some interest to examine the
circumstances that make the ESO
domes visible over such distances. Key
clues are the time at which they may be
seen, from about 15 to 10 minutes be-
fore apparent sunset at GTIO, and the
4-5-minute period over which they are
seen. Such a short timespan suggests a
Specular or semi-specular reflection off
a surface which is flat in at least one
dimension. Reflections from the domes
themselves are thus ruled out, since the
brightness of such a reflection would
Change very little over most of the day.
Reflections from the flat walls of struc-
tures such as the Administration Build-
ing would be prominent indeed, but
could not be seen over a range of solar
declination of more than 0.5 degree, in
Conflict with the observed month-Iong
range of visibility. The cylindrical sides
of many of the domes are the most likely
source of short-lived reflections, but ad-
ditional checks are needed to make sure
their properties are fully consistent with
the time of appearance and duration of
the observed reflections.
Since GTIO and ESO are at compar-
able elevations, to a good approxima-
tion the condition for reflection off a
Vertical surface is that the sun should
have an apparent zenith distance of 90
degrees at La Silla. This astronomical
condition can be tested from the known
Situation of the observatories and the
known time of visibility relative to t11e
apparent sunset at Tololo. First, how
long before apparent sunset at La Silla
is the apparent sun horizontal, and sec-
ond how much difference in the times of
sunset is there between La Silla and
The dip of the horizon from the al-
titude of ESO is about 1.6 degrees, and
at this time of year the sun requires
about 435 seconds to traverse this verti-
cal distance solely from geometrical
considerations. An additional time be-
tween the sun at 90 degrees zenith dis-
tance and apparent sunset is produced
by atmospheric refraction, which pro-
duces a deceleration of the apparent
sun amounting to about 160 seconds for
typical meteorological conditions; most
of the refraction seen at sunset arises in
the last few degrees above the horizon.
Finally, apparent sunset at Gerro Tololo
is 80 seconds later than at La Silla due
to the smalilongitude difference. Adding
these up, the horizontal-sun condition is
satisfied from about 13 to 11 minutes
before apparent sunset on Gerro Tololo,
in good agreement with what we actual-
Iy see. The reflections are bright for
somewhat longer than the 2 minutes
expected from the sun's angular size
and a perfectly smooth surface, pre-
sumably due to the roughness of the
corrugated metal used for the cylindrical
sections of some of the telescope build-
ings. (Unfortunately for sunset watchers
on La Silla, most of the domes at Gerro
Tololo have rectangular base structures
and thus any reflections will be visible
for only a couple of days each year;
furthermore they are at unfavourable
orientations for catching sunlight as
seen from almost due north).
Examining the photograph once
more, the 3.6-m dome is the brightest
(far right), as befits the large size of the
support building. One can also identify
(Ieft to rigllt) the small cluster of domes
dominated by the Dutch 90-cm, the
ESO 1.5-m, the ESO 1.0-m and Danish
1.5-m (both rather faint), and the ESO/
MPI 2.2-m. It is fitting that visibility
phenomena of this kind may be under-
stood with some of the astronomical
observer's most basic methods.
FERRARO, Francesco (I), Fellow
KÄUFL, Hans Ulrich (0), Infrared
Instrumentation Scientist
PASIAN, Fabio (I). Fellow (Senior Arcllive
Scientist ST-ECF)
BUYTENDIJK, Felice (NL), Receptionist
GROTE, Rainer (CH), Projekt
MEYLAN, Georges (CH), Fellow
SCHNEIDER, Karin (0), Secretary
VAN RIJN, Gunilla (NL), Administrative
HAGSTRÖM, Magne (S), Associate
(Microwave Engineer SEST)
OLBERG, Michael (0), Telescope
Software scientist
Signposts 01 Low Mass Star Formation in Molecular Clouds
B. REIPURTHand C. MADSEN, ESO
Regions of massive star formation are
easily recognizable because of the pres-
ence of bright, often very extended H11
regions. The more quiescent places
where only low mass star formation
takes place are not so immediately obvi-
ous to identify. Most of the low mass
star forming regions known today were
found in the 1950's, mainly through ob-
jective prism surveys for Ha emission
stars done by Joy, Herbig, Haro and
others. Their results are summarized
and supplemented by later findings in a
new catalogue by Herbig and Bell
(1988), which Iists 742 mainly low mass
pre-main-sequence stars. Another re-
cent and rich source of low mass young
stars is the IRAS catalogue. IRAS data
towards clouds, however, often suffer
from source confusion and, in particular,
extraction problems because of
background emission.
Reflection and Emission Nebulae
The youngest and generally most in-
teresting stars are still intimately associ-
ated with their placental material, and
therefore often show smalI, faint reflec-
ti on nebulae. In the earliest evolutionary
stages a star is normally not visible at
all, but through cavities blown open to
the cloud surface it may illuminate the
surrounding cloud. Perhaps the finest
known case is the Re 50 nebula in Orion,
shown in a CCO image through a red
broadband filter in Figure 1. This object
was found by inspection of deep red
Schmidt plates of the L 1641 molecular
cloud in Orion (Reipurth 1985). The
northern elongated nebula is a beam of
light that escapes from a 250 ~ em-
bedded infrared source. The southern
nebula is a molecular cloud clump
which is iIIuminated by the same star
through a channel hidden from view.
The whole object is highly variable, to
the extent that it looks different each
time it has been observed during the last
6 years. Moreover, inspection of older
Palomar Schmidt plates from 1955
shows nothing of the nebula. Perhaps
we are here for the first time witnessing
the emergence of a young star from its
cloud-enshrouded birthplace (Reipurth
and Bally 1986, Scarrot and Wolsten-
croft 1988). Some other similar reflec-
tion nebulae, apparently in slightly more
evolved stages, are the PV Cep, Re 5
and IRN nebulae (Cohen et al. 1977,
Graham 1986, Schwartz and Henize
After a young star has broken through
its cloud cover to the outside world it
passes through phases where strong
outflowing winds create shocks in the
ambient medium, the so-called Herbig-
Haro objects. These are normally found
as chains of small nebulae stretching
away from an infrared source. They can
be identified by their very characteristic
emission line spectra. In recent years it
has become clear that a subset of the
Herbig-Haro objects take the form of
highly collimated jets, remarkably simi-
lar in appearance to the jets from ex-
tragalactic radio sources (for arecent
review, see Mundt 1988).
One of the finest jets known to be
associated with a young star is the HH
34 jet. It was originally noted on a deep
red Schmidt plate, and subsequent
CCO imaging revealed its extraordinary
nature (Reipurth 1985, Reipurth et al.
1986). Figure 2 shows an almost 30 arc-
second long jet emerging from a faint
emission line star, and pointing right to-
wards HH 34 in the lower part of the
image. We are here witnessing mass
loss from the star ejected in a collimated
supersonic beam, ending in a bow
shock where the outflowing material
rams into the ambient medium. A com-
parison between the original discovery
image from 1982 with aseries of identi-
cal images taken up to 1989 shows that
the shock structures seen as knots in
the jet have a proper motion in the direc-
tion away from the source. A second
oppositely oriented bow shock was
found on the other side of the jet by
Bührke et al. (1988), demonstrating that
the outflow is bipolar. Studies of such
objects provide key information on the
earliest phases of stellar evolution.
Giyen the success of the initial exami-
nations of Schmidt plates in identifying
new and unusual regions of low mass
star formation, a more ambitious project
was initiated.
A Survey of Molecular Clouds
with the ESO Schmidt Telescope
Because of their large fjeld, Schmidt
plates are ideal for surveys. The fine-
grained 111 a-F plates are, when properly
hypersensitized, particularly useful for
searches for the often intrinsically red
and heavily obscured small reflection
nebulae, as weil as for Herbig-Haro ob-
jects, since their emission lines mainly
fall within the spectral sensitivity curve
of the 111 a-F emulsion. Using the ESO
1-m Schmidt telescope at La Silla, Chile,
equipped with a RG 630 filter, a large
scale survey of molecular clouds all
along the southern galactic plane has
been carried out during the last two
years. All plates were taken by H.-E.
Schuster, Guido Pizarro and Oscar
Pizarro. More than two hundred ti ny
nebulae in dark clouds have been iden-
Seen on deep Schmidt plates, ex-
posed to their best S/N ratio, such small
nebulae appear to be superimposed on
a background of relatively high density.
Thus at times it is difficult to recognize
the tiny objects by merely visually in-
specting the plates. The weil known
photographic technique of diffuse-light
amplification brings remedy to the prob-
lem by "removing" the chemical fog
contributing to the high overall density,
and by increasing the contrast of the
image at the density level required.
However, as such nebulae are some-
times found in regions with nearby OB
stars surrounded by bright H11 regions,
there can be large density variations
over the photographic plate. Therefore,
pure amplification will effectively ob-
scure as much as it reveals, due to the
very Iimited dynamic range which this
process offers. Consequently it is often
necessary to employ a strong unsharp
mask during the amplification process.
The technique of unsharp masking has
been described elsewhere and shall not
be dealt with here. Serving as an ana-
logue low frequency filter, the mask
brings down the overall contrast of (a
larger part of) the plate to a level which
allows for contrast amplification of the
whole area under study.
The follow-up Observations
In order to separate reflection from
emission nebulae, one should ideally
obtain a spectrum of each nebula. How-
ever, most of these tiny objects are so
faint that it would require immense
amounts of observing time on large tele-
scopes to complete such a programme.
But another simple and fast technique
exists. For each object direct CCO im-
ages were obtained at the Oanish 1.5-m
telescope at La Silla, one through a far-
red broadband Gunn z filter extending
to the CCO cutoff beyond 1 micron, and
one through a narrow-band interference
filter centred on the [S 11] 6717/6731
emission Iines. A reflection nebula
associated with a partly embedded
young low-mass star is very red and will
show up prominently in the broadband
Gunn z filter, but rarely in the narrow-
band sulphur filter. On the other hand,
Herbig-Haro objects are strong emitters
in the 6717/6731 lines, but have no
Figure 1: A Gunn r filter GGO image taken 4 March 1986 at the Oanish
1.S-m telescope. An infrared source is hidden above the upper
nebula, which is a beam of light escaping from the newborn star. The
lower nebula is a molecular cloud clump illuminated by the embed-
ded star.
Figure 2: A [Sill filterGGO image taken at the Oanish 1.S-m telescope
of the HH 34 region. Here a shocked emission line jet flows super-
sonically away from a faint young star. The jet cools to invisibility, but
continues and eventually rams into the ambient medium in a bow
strong Iines in the Gunn z bandpass and
Virtually no continuum, so they will show
up very bright in the narrow-band sul-
Phur filter, and barely, if at all, in the
Gunn z filter. In this way the number of
known Herbig-Haro objects have been
doubled. Since such objects have most
of their emission concentrated in a cou-
pie of strong emission lines, it is feasible
to study spectroscopically this subset of
the list of new small nebulae. These
additional observations are important
for final confirmation, since distant H11
regions sometimes can mimic HH ob-
jects. But the gradual build-up of experi-
ence hel ps: at the beginning of the sur-
V f ~ y only 15 % of the objects which on
the plate looked to be Herbig-Haro ob-
jects were actually confirmed as such,
while at the end the rate exceeded 90 %.
It should be noted that the selection
of a [SII] 6717/6731 filter, rather than an
Ha filter, is essential to this technique,
as most stars show Ha emission, which
will then also be present in the
Surrounding reflection nebula. [S 11]
emission is, on the other hand, a clear
indicator of shocks, when associated
with a low-mass, low-Iuminosity star.
Van den Bergh (1975) employed an early
version of this scheme.
Such a survey has a value in providing
firmer statistics on the frequency and
timescales of shocked outflows from
young stars. But it is the detailed follow-
up observations of selected objects
which provide the most fun and which
occasionally can provide new insights
on early stellar evolution. For example,
the new objects HH 80/81 are almost as
bright as HH 1 and 2 (the first objects to
be discovered by Herbig and Haro and
still the brightest), but they are at least
three times more distant and are thus
the intrinsically brightest objects known
(Reipurth and Graham 1988). Moreover,
they do not emerge from a solar-type
star, but from a young B-star, and they
show the highest velocity dispersion
hitherto observed. Or take the case of
the new object HH 111, probably the
largest and best collimated jet known to
date to emerge from a young star. Be-
sides the jet it has four bow shocks,
moving supersonically in pairs of two in
opposite directions from an embedded
infrared source, from which there also
streams a large molecular flow
(Reipurth, in press).
The Orion Region
The molecular clouds in Orion are the
most active sites of low mass star for-
mation known. Figure 3 shows a deep
red Schmidt plate of the Orion nebula
and beneath it the L 1641 molecular
c10ud stretching to the southeast. The
whole region is full of many hundreds of
Ha emission stars, variable stars, flare
stars and infrared sources. The L 1641
Figure 3: The cometary shaped molecular
cloud L 1622 in Orion as seen on a deep red
ESO Schmidt plate. The cloud is active in
low-mass star formation, and contains many
nebulous stars, Ha emission stars, infrared
sources and one Herbig-Haro object. ~
t :'
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Figure 4: Adeep, red ESO Schmidt plate (lIIa-F, 150 min., RG 630) showing the Orion Nebula. The print was produced by subjecting the plate to
unsharp masking, followed by diffuse-light amplification. As M42 itself is not interesting in this context, the mask was prepared in such a way
that it would not influence M42 in any appreciable manner, but so that the low-surface-brightness filamentary structures stand out in great
cloud appears as a rather homogeneous
obscuration, but a detailed map made at
13CO by Bally et al. (1987) shows that
the cloud is really composed of long
infiltrated strings and clumps of denser
material embedded in a lower density
gas and dust environment. Closer ex-
amination of the region shows that the
youngest stars, those found as infrared
sources, are almost invariably located in
association with the denser clumps.
The region is also rich in Herbig-Haro
objects, indeed it has the highest con-
centration of such objects found any-
where. Star formation therefore has not
only taken place here over the last few
million years, but is a still ongoing pro-
cess. Moreover, it is not confined merely
to the dense L 1641 cloud. Figure 3 has
been processed in such a way as to
bring out the faint outlying structures, to
which only little attention has been paid
to date. Our survey uncovered several
Herbig-Haro objects in these more
peripherical regions, and shows that
star formation also occurs here, albeit
on a very modest scale.
Orion contains many smaller, less
well-known clouds, in which stars are
being born. Figure 4 (centerfold) shows
the beautiful L 1622 c1oud. With its
bright rims and long tail it gives the
impression of being strongly affected
from the outside. It is in fact pointing
directly towards the young massive 0-
stars in M42, which bath it in ultraviolet
Professors J. H. Oort, H. van der Laan and A. Blaauw looking at the La Silla model.
radiation. This outside influence may be
the cause of the vigorous formation of
stars occurring in this small cloud: there
are several nebulous stars, Ha emission
stars, and embedded infrared sources,
and we have found a new Herbig-Haro
object (HH 122, seen as a tiny group of
small nebulae near the eastern edge of
the cloud). A detailed optical infrared/
radio study is currently being made of
L 1622 at La Silla.
We are very grateful to Hans-Emil
Schuster, Guido Pizarro and Oscar
Pizarro, who took all the plates used in
this survey.
Bally, J., Langer, W.D., Stark, AA., Wilson,
R. W.: 1987, Astrophys. J. 312, L45.
Bührke, T., Mundt, R., Ray, T. P.: 1988, As-
tron. Astrophys. 200, 99.
Cohen, M., Kuhi, L.V., Harlan, E.A: 1977,
Astrophys. J. 215, L127.
Graham, J.A: 1986, Astrophys. J. 302, 352.
Herbig, G. H., Bell, K. R: 1988, Liek Observa-
tory Bulletin No. 1111.
Mundt, R: 1988, in Formation and Evolution
of Low Mass Stars, eds. A K. Dupree,
M.T.V.T. Lago, NATO ASI Series C vol.
241, p. 257.
Reipurth, B.: 1985, Astron. As/rophys. Suppl.
Reipurth, B., Bally, J.: 1986, Nature 320,336.
Reipurth, B., Bally, J., Graham, J.A, Lane,
A.P., Zealey, W.J.: 1986, As/ron. Astro-
phys. 164,51.
Reipurth, B., Graham, J.A: 1988, Astron. As-
trophys. 202, 219.
Scarrott, S. M., Wolstencroft, R 0.: 1988,
Monthly Notiees R. Astron. Soe. 231, 1019.
Schwartz, R.D., Henize, K. G.: 1983, Astro-
phys. J. 88, 1665.
van den Bergh, S.: 1975, Publ. Astron. Soe.
Paeifie 87,405.
On Friday, January 27, Professor
Adriaan Blaauw opened the ESO Ex-
hibition in The Hague, the Nether-
lands, with a review of ESO's history
since 1952. This was preceded by a
speech from the Director General
about ESO's future and its role in
European astronomy.
Among many prominent guests
were Prof. Jan Hendrik Oort and Dr.
Henk Bannier, both former Presidents
of the ESO Council, and the present
members of Council, Prof. Wim Brouw
and Dr. Jan Bezemer.
The Exhibit, which lasts till March
12, 1989, is hosted by the beautiful
new Science Museum of The Hague
called MUSEON. ESO is especially
grateful to Dr. Wim van der Weiden,
MUSEON's Director, for his enthusias-
tic reception. The Exhibit was set up
by Mr. Claus Madsen of ESO and Dr.
Peter Wisse, staff astronomer of MU-
The festive opening was co-hosted
by OMNIVERSUM, Europe's first
space-theatre, next door to MUSEON.
OMNIVERSUM openend in December
1984 and is the result of a sustained
initiative by Prof. Harry van der Laan
between 1977 and 1984, while he was
Chairman of nearby Leiden Observa-
Observation ofthe 12CO (J -1 ~ 0) Line in NGC 613
with the SEST
E. BAJAJA*, Max-Planck-Institut für Radioastronomie, Bonn, F.R. Germany
E. HUMMEL, University of Manchester, Nuffield Radio Astronomy Laboratories, Jodrell Bank,
Macclesfield, Cheshire, U. K.
The availability of the Swedish-ESO
SUbmillimeter Telescope (SEST) on La
Silla opened the possibility of extending
the radio observation of molecular lines
to the very southern galaxies. In particu-
• Member of lhe Carrera dei Invesligador Cientlfico
of lhe CONICET, Argenlina.
lar the observation of the CO lines in the
nearest galaxies will permit not only to
increase the sampling for statistical pur-
poses but also to study in more detail
the distribution and kinematical proper-
ties of the molecular clouds in relation to
other components of the galaxies. The
HPBW of the SEST, at the frequency of
115 GHz of the 12CO (J = 1 ~ 0) line, is
43" which means that galaxies with
diameters between 5.5 and 10 minutes
of arc are weil suited for mapping since
they do not require a prohibitive amount
of time and the arms can be resolved if
the inclination angle is adequate.
We had selected NGC 613 some time
aga as a candidate for CO observation
because of several interesting features,
Figure 1: The 12CO (J = 1 -40) profiles (veloeity range 1,000 to 2,000 km s-', resolution 7.2 km S-I) in NGC 613 overlayed on an optieal pieture of
the galaxy. Coordinates on the frame indieate offsets with respeet to the eentre. North is at the top, east to the left.
which will be described in the following
section, and since it fulfills the size con-
dition (diameter about 6') we proposed
it for the observation with the SEST. The
observations were made in July 1988.
Here we present the preliminary results.
NGC 613 is a moderately southern
(0 '"' -29<:> 66) barred spiral galaxy,
classified as SB(rs) bc by de Vau-
couleurs (1976). The bar and the several
arms are weil delineated by prominent
dust lanes and H11 regions. The galaxy
has been studied with optical spectros-
copy by Burbidge et al. (1964) who were
able to determine the position angle
(115°), the central velocity (1,534 km/s,
heliocentric) and the rotation curve.
It has also been observed in the radio
continuum at ), 6 cm and /... 20 cm and in
the Ha and [0111] ),5007 lines by
Hummel et al. (1987). These observa-
tions show that there are jet-like struc-
tures extending from the centre indicat-
ing the presence of nuclear activity. The
linear scale of these features is of the
order of 5" to 30", thus unresolvable by
the SEST beam.
The galaxy has been observed previ-
ously in the 12CO (J = 1 --70) line by
Elmegreen and Elmegreen (1982) who
reported a marginal detection at the
centre with an upper limit of 0.04 K us-
ing the NRAO 10m dish at Kitt Peak.
Observations in the HI 21 cm line were
made by Bajaja (1978) and Reif et al.
Observation and Reduction
The SEST has been described by
Booth et al. (1987). The antenna para-
meters, at /...2.6 mm, are: HPBW = 43",
beam efficiency = 0.78, aperture effi-
ciency = 0.67. At the front-end a dual
polarization 85-117 GHz receiver with
cooled Schottky diode mixers and a sin-
gle side-band noise temperature of
250 K was available. At the back-end
we used a wide band Acousto-Optic
Spectrometer (AOS) with 1,700
channels spaced 690 kHz (1.8 km S-l at
/...2.6 mm).
The observations were made between
the 25th and the 29th of July, 1988. The
beam was switched in position, every 2
minutes, between the source and a ref-
erence point at 12' to the west. The
integration time per point was, in gener-
al, 16 minutes but some positions,
especially along the major axis, were
observed longer. A total of 27 points
were observed on a square grid with 20"
spacing, aligned with the major and
minor axis (position angle = 115°), cov-
ering the central part of the galaxy. Dur-
ing the time these observations were
being made, it was realized that there
was a problem with the telescope point-
ing. Sudden jumps of the order of l'
occurred near azimuth = 0°. Since
NGC 613, however, was observed from
rising to meridian crossing, only few
spectra, which could be easily iden-
tified, were lost because pointing para-
meters used by previous observers
were adopted. We believe, that the
accuracy of the positions is within 5"
which is 12 % of the beam size.
The spectra were saved on tape with
FITS format and read at the MPlfR Mi-
crovax converting it to CLASS readable
format. This facility was employed to
correct the baselines and to do Hanning
smoothing. The average rms noise, in
the spectra with a velocity resolution of
3.6 km/s, is 0.03 K of corrected antenna
Booth, R.S., Oe Jonge, M.J., Shaver, P.A.:
1987, The Messenger48, 2.
Burbidge, E.M., Surbidge, G.R., Rubin, V.C.,
Prendergast, K. H.: 1964, Astrophys. J.
Elmegreen, D.M., Elmegreen, B.G.: 1982,
Astron. J. 87, 626.
Hummel, E., Jörsäter, S., Lindblad, P.O.,
Sandquist, A.: 1987, Astron. Astrophys.
Reif, K., Mebold, U., Goss, W.M., van Woer-
den, H., Siegman, S.: 1982, Astron. Astro-
phys. Suppl. Sero 50, 451.
Vaucouleurs, G. de, Vaucouleurs, A. de, Cor-
win, H.C.: 1976, Second Reference Cata-
logue of Bright Galaxies, The University of
Texas Press, Austin.
1000 1200 1400 1600 1800 2000
Figure 2: Position-velocity diagrams along (a)
the major axis, (b) the minor axis. Velocities
are heliocentric. Ordinates indicate offsets,
along each axis, with respect to the centre of
the galaxy. First contour lines (dashed) corre-
spond to the level 0.02 K of antenna temper-
ature. The contour interval is 0.02 K.
Bajaja, E.: 1978, Publ. Dept. Astron. Univer-
sidad de Chile, 111, 55.
We are grateful to all the ESO staff
members who made these observations
minor axis (Fig. 2 a) the velocities cover
a range of about 275 km S-1. This is a
consequence of the steepness of the
rotation curve and the width of the
beam. The CO emission is asymmetri-
cally distributed, being mainly concen-
trated on the SW side.
The derivation of the molecular gas
mass from the areas of the profiles of
Figure 1 depends on the distance to the
galaxy and on the conversion factor to
column densities. Both contain large
uncertainties. Assuming a distance of
19 Mpc as derived from the systemic
velocity, corrected for the motion with
respect to the Local Group, assuming a
Hubble constant of 75 km S-1 Mpc-
and a conversion factor of 410
(K km s-1r
, the mass of the gas in the
form of molecular hydrogen would be
about 3.1 10
MG' The mass of the neu-
tral hydrogen estimated by Reif et al.
(1982) is about 3.6 10
MG' which
means that the total gas (HI + H
) mass
would be about 6.7 10
MG' Assuming
an inclination angle of 45 degrees, the
highest rotational velocities, which are
measured at about 80" from the centre,
(Fig. 2b) are of about 212 km/so With
these values it is possible to estimate
roughly the total mass within that radius
as about 7.610
MG' The ratio between
the gas and the "total" mass would then
be of the order of 9 % if the assumed
values for the parameters used in these
calculations are valid.
Figure 1 shows the obtained profiles.
The CO profiles were Hanning
smoothed two times so the velocity res-
olution of the displayed profiles is
7.2 km S-1 and the rms noise goes from
0.01 K to 0.03 K depending on the in-
tegration time and observing conditions.
The velocity range on each profile is
1,000 to 2,000 km S-1 from left to right.
Along the major axis we have the pro-
files with the lower noise. In spite of the
fact that one profile is missing in the
sequence (at the offset 60" to the SE)
the position-velocity diagram (Fig. 2b)
clearly shows the velocity variation
along this axis. From this diagram it is
possible to determine the heliocentric
velocity of the centre of the galaxy
which is 1,470 ± 10 km S-1 and the rota-
tion curve (correcting for the inclination).
The differences between this central
velocity and the velocities obtained from
global HI profiles by Bajaja (1978)
(1,431 km S-1) and by Reif et al. (1982)
(1,493 km S-1), are consistent with the
noise in the spectra. The difference with
the velocity derived by Burbidge et al.
(1964) from optical observations, how-
ever, is rather large (64 km S-1). The
highest velocities along the line of sight,
with respect to the centre, in Figure 2 b,
are about 150 km S-1. This value is also
smaller than the highest velocities seen
optically by Burbidge et al. (200 to
250 km S-1). At both ends of the dia-
gram appear features at velocities that
correspond to the other side of the
galaxy which would imply rotation in the
opposite sense. These features are
most probably due to the noise. It
should be mentioned that Burbidge et
al. (1964) also found strong irregularities
at both sides of the centre, although
neither at the same distance nor with the
same velocity differences. Along the
High Resolution Ha Spectroscopy of Nova Centauri 1986:
Tracing a Transient in the Spectral Evolution
R. BAND/ERA, Osservatorio Astrofisico di Arcetri, Firenze, /ta/ia
P. FOCARD/, Dipartimento di Astronomia, Universita di B%gna, /ta/ia
We present here some preliminary re-
sults and considerations regarding high
resolution spectroscopy of Nova Cen-
tauri 1986. This object was discovered
on November 22.7 UT as a 5.6 V mag-
nitude star (1), and observed by us in Ha
two months later, when it was of about
10 magnitude. A considerable change in
the Ha line profile, basically consisting
in the disappearance of the original f1at-
top profile substituted by a more regular
multicomponent profile, is clearly visible
in our set of spectra, obtained monitor-
ing the nova for nearly a week. A fitting
of the line profile by means of three
gaussian components shows a narrow-
ing of the components and a rapid fad-
ing of one of them as a function of time.
This behaviour can be interpreted in
terms of blobs which have been formed
in a non-spherical explosion and that
6550 6560 6570 6550 6560 6570 6550 6560 6570
Figure 2: Comparison between Spectra # 2-5 (increasing number, i. e. increasing time,
downwards). The original data are on the left; the fitted models, in the middle, are composed of
3 gaussians plus a second-order continuum; the residua/s after fitting are shown on the right.
All spectra are normalized; dashed fines give the zero levels of spectra. The dates of the
observations (see Table 1) are reported on the lett side.
22.36 .
6540 6550 6560 6570 6580
Figure 1: Comparison between Spectrum
# 1 (upper spectrum) and Spectrum # 2
(Iower spectrum). 80th spectra are nor-
mafized; dashed fines give the zero levels for
the two spectra. The dates of the observa-
tions (see Table 1) are reported on the left
The relative amplitude of components
1 and 2, in the main peak, remain stable
with time, while component 3 fades with
respect to the others with a time scale of
5 days (computed on the last 3 spectra).
All components present a slight de-
crease of their width, as weil as a slight
shift of their central wavelengths with
time. The continuum cannot be deter-
mined precisely: its slope seems to de-
crease with time, and the data are con-
sistent with an intensity evolving as that
of the former two components.
Is there any physical meaning for the
three gaussian components fitted
above? One could alternatively imagine
that an otherwise symmetric Une is par-
tially absorbed on its blue wing, by a P
Cygni-Iike effect. However, one would
expect this effect to disappear with
time, while the observed trend is the
opposite. Furthermore, by fitting with
gaussians the positive peak of a P Cygni
line, the central wavelength of these fic-
tious components will shift redwards
when the effect is stronger; the derived
parameters, instead, go in the opposite
way. Therefore we think that the line is
physically composed of emission com-
ponents. What we fitted with the former
two components could actually be a
Spec- Date Exposure
trum (Ul) Time (s)
No. 1 Jan. 13.36 1200+900
NO.2 Jan.17.34 600+600+600
No. 3 Jan. 19.35 600+600
NO.4 Jan.20.33 600+600
No. 5 Jan.22.36 600+600
mated by A + B (I.. -6560) + C (I..
This fit gives good results: synthe-
sized profiles (Fig. 2, middle) reproduce
the original ones weil, and the residuals
are limited to a few per cent (Fig. 2, right
side). The main parameters of the spect-
ra are listed in Table 2. The time de-
pendence of the various parameters is
rather clear, with the exception, for a
few of them, of the Jan. 17.34 spectrum:
this probably because in this spectrum
the components are too blended to be
separated correctly, and therefore the
parameters derived from it are less
The Evolution of Spectra
Spectrum # 1 appears as a flat-top
spectrum, with many subcomponents
and a red wing, in reasonable agree-
ment with high-resolution spectra ob-
tained by Pacheco (private communica-
tion) during the period 4-9 January, us-
ing the 1.4-m CAT telescope equipped
with a Reticon detector. With respect to
Pacheco's spectra, the Ha Une is
narrower, the red wing is less evident,
while a blue component becomes vis-
After only 4 days (spectrum # 2) the
Ha profile has changed considerably
(see Fig.1), taking on a more regular
shape. Further variations are seen by a
direct comparison of spectra # 2-5
(see Fig. 2, left side): the Ha profile is
getting thinner and thinner; in the
meanwhile, the component on the blue
side is getting fainter and fainter.
Derived Parameters
To give a quantitative description of
the spectral evolution we fitted the pro-
files # 2-5 with three gaussians (two
for the main component, plus one for
the blue component), added to a sec-
ond-order continuum: the gaussian
components are labelled from 1 to 3
bluewards; for each of them Ai indicates
the line amplitude, Ai the central wave-
length, and W
the full width half maxi-
mum, while the continuum is approxi-
were evaporating at the time of the ob-
The Observations
The spectra were obtained during the
"spare time" of an observing run with
the 1.4-m CAT telescope, and the CES
spectrograph equipped with the Short
Camera (2) and a double-density
RCA CCO. Our main goal was a search
for binarity in P Cygni-type stars and red
supergiants: part of the results we ob-
tained can be found in 3, 4. The instru-
mental setting is of extremely high effi-
ciency, allowing to obtain spectra of
superb quality with a S/N ratio exceed-
ing 200. The slit width was 2 arcsec on
the sky, corresponding to a resolving
power of nearly 40,000; the original
spectra were = 50 Awide, centred on
the wavelength 6560 A. The log of ob-
servations is reported in Table 1. Pre-
liminary reduction (flat field correction,
wavelength calibration) has been per-
formed with the MIOAS package avail-
able at the Bologna and Arcetri In-
stitutes. For these spectra, no absolute
measurement is available to us; there-
fore they will be presented here, nor-
malized to their peak value.
Date: Jan.17.34 Jan. 19.35 Jan.20.33 Jan.22.36
~ , 6566.0 6565.4 6565.3 6564.7
~ 2 6560.8 6560.9 6560.9 6560.5
~ 3 6553.4 6554.0 6553.4 6552.5
6.6 5.3 5.2 5.0
8.7 9.7 9.2 8.6
1.01 0.98 1.06 1.04
7.1 10-
unique non-gaussian component: how-
ever, the result of the two-component fit
looks remarkably good.
It is difficult to explain the details of
this line evolution by a spherically
symmetrical model and, on the other
hand, there is evidence for departures,
in nova events, from spherical symmetry
(e. g. 5, 6). In our case one could imagine
that, while most of the nova shell dis-
appeared due to dilution effects, some
denser blobs began dominating the Ha
emission. The velocity dispersions of
the three components are respectively
190, 140 and 250 km/s, and represent
their expansion velocities.
Assuming that all blobs have similar
sizes, the time scale for the fading of the
li ne emission is inversely proportional to
the blob expansion velocity, and in fact
component # 3, the broader one, gets
dimmer faster than the others. Let us
define 1:12 as the time scale for the fading
of the former two components, and 1:3 as
that for the latter one; since the de-
crease of component # 3 relative to the
others is 5 days, its absolute time scale
is 1:3 = 1/(1/5 + 11t
). Requiring also that
time scales are inversely proportional to
the components' expansion velocities,
1:12 and 1:3 can be derived separately as,
respectively, 3 and 1.9 day.
Unfortunately we da not possess ab-
solute calibration of our spectra; any-
how, we observed that the continuum
was evolving, in those days, as the com-
ponents # 1 and # 2. But a rough esti-
mate of the continuum evolution can be
given by taking the decrease of the visu-
al magnitude: by using visual mag-
nitudes measured during the month of
January (7, 8), the average decrease
turns out to be 0.35 mag/day, corre-
sponding to a time scale of 3.1 day, in
e10se agreement with our previous deri-
A decrease in line amplitude can then
be ascribed to expansion of condensa-
tions in the nova envelope; they have
typically a size of "" 10
cm, much
smaller than the envelope size at that
time ("" 5.10
cm); since the compo-
nents are partially blended, they should
be confined to the inner regions
("" 10
Therefore a way to explain the ob-
served Ha behaviour of Nova Centauri
1986 involves the presence of a non
spherical explosion, with the formation
of small condensations.
We are indebted to M. Friedjung and
E. Oliva for interesting suggestions, and
to J.A. de Freitas Pacheco for giving us
a copy of his spectra.
(1) McNaught, RH. 1986, lAU Gire. N. 4274.
(2) Dekker, H., Delabre, B., D'Odorico, S.,
Lindgren, H., Maaswinkel, F., and Reiss,
R 1986, The Messenger43, 27.
(3) Gratton, RG., Focardi, P., and Bandiera,
R. 1989, M.N.RAS., submitted.
(4) Bandiera, R, Focardi, P., Altamore, A.,
Rossi, C. and Stahl, O. 1988 Physies of
Luminous Blue Variables, lAU Coll.
No. 113, in press.
(5) Mustel, E. R, and Boyarchuk, A.A. 1970,
Ap.Sp.Sei.6, 183.
(6) Hutchings, J.B. 1972, M.N.R.A.S. 158,
(7) McNaught, RH., Campos, J. 1987, lAU
eire. N. 4298.
(8) McNaught, RH., 1987, lAU Gire. N. 4315.
The European Working Group on Chemically Peculiar Stars
of the Upper Main Sequence: The First 10 Years
G. MATHYS (Geneve), H. M. MAITZEN(Vienna), P. NORTH(Lausanne),
H. HENSBERGE(Brussels), W. W. WEISS (Vienna), S. ANSARI (Vienna),
F.A. CATALANO (Catania), P. DIDELON(Strasbourg), R. FARAGGIANA (Trieste),
K. FUHRMANN (Göttingen), M. GERBALDI (Paris), P. RENSON(Liege),
H. SCHNEIDER(Göttingen)
1. Introduction
A fraction of the Band A type stars
have chemical peculiarities (CP). 4 sub-
groups are recognized: CP1 or Am,
CP2 or magnetic Ap (with enhanced Sr,
er, Eu, Si lines), CP3 (ar non-magnetic
Ap, with Hg Mn enhanced) and CP4 (B
type stars with He peculiarities, a frac-
tion of them appear to be a hot exten-
sion of CP 2).
For 10 years now, European as-
tronomers interested in the study of CP
stars, and more particularly of Ap-Bp (or
CP2 to CP4) stars, have gathered their
efforts in a working group (WG). Many of
the results obtained by members of this
WG have been derived from observa-
tions obtained at ESO-La Silla. In the
Messenger No. 34, an overview of the
activity of the group during its first five
years of existence was given. On the
occasion of the 10th anniversary of our
WG we want to present a new report on
its work with emphasis on those studies
carried out during the last five years.
Beside the original research work of
the various group members, a Catalogue
of CP-stars has been compiled by P.
Renson. This list was set up after a large
critical survey of the literature; it con-
tains more than 6,000 stars, half of
which are CP1 (i. e. metallic line) stars. A
number of remarks are provided for
each star. They concern mainly other
designations, misidentifications, duplici-
ty, variability of different parameters
(light, spectrum, magnetic fjeld). With
this catalogue, which will soon become
available through the "Centre des
Oonnees Stellaires" (Strasbourg), will
also be published a list of bibliographi-
cal references established by P. Ren-
son, M. Gerbaldi, F. Catalano and
M. Floquet.
2. Variability
Variability in Ap and Bp stars occurs
in (at least) two flavours: on the one
hand with periods ranging fram the or-
der of half a day to several years, on the
other hand with short periods of about 4
to 15 minutes. These latter, rapid varia-
tions will be discussed in a following
section. The former type of variability is
ubiquitous among CP2 and CP4 stars,
almost all of which undergo periodic
variations of their brightness, spectral
features and magnetic field. For most of
the cases these are weil explained by
the Oblique Rotator Model: the star is
regarded as a rigidly ratating body with
an essentially dipolar magnetic field,
whose axis does not generally coincide
with the rotation axis. Surface in-
homogeneities likely related to the
magnetic structure are responsible for
the observed variations, as a result of
the changing aspect of the visible stellar
hemisphere along the rotational cycle.
However, as will be argued below, it is
not yet quite clear whether all the ob-
served "slow" (as opposed to rapid) var-
iations of CP stars can be explained in
terms of the Oblique Rotator Model. In
particular, some CP3 stars appear to
show photometric variability (studied at
ESO mainly by Schneider, 1987) with
amplitudes generally lower than for CP2
and CP4 stars.
The most convenient way to deter-
mine the rotational period is to make
precise, multicolour photometrie obser-
vations. The Strömgren system has
been widely used for this purpose, and
Geneva photometry has contributed to
this effort.
All periods known up to the beginning
of 1983 have been gathered in a cata-
logue published by Catalano and Ren-
san (1984). About 250 stars had a
known period at that time. The first sup-
plement has just been published
(Catalano and Renson, 1988) and con-
tains further 58 stars, so that more than
10 per cent of all known Ap stars have a
known rotational period.
The contribution of the members of
the European WG is quite important,
since three quarters of the new periods
published since 1983 come fram them.
Of these (representing 41 stars) most of
them have been published by North
(1984: 5 stars, 1987: 18 stars) and
Manfraid and Mathys (1985: 5 stars,
1986: 8 stars). In these statistics we did
not include the few CP eclipsing
binaries, but added 3 stars not con-
tained in Catalano and Renson (1988).
The new periods found since 1983 are
presented in the histogram of Figure 1.
Compared to the histogram shown in
the Messenger No. 34, there is an ex-
cess of short periods. The scarce
appearance of periods shorter than 0.5
days is worth mentioning; only one of
them is quite certain. The other three
should still be confirmed.
Now, what is the "completeness" of
the period determinations? In order to
answer this question, we selected all
Southern (i. e. with negative declina-
tions) Ap stars with V < 7.1 (which are
members of the BS catalogue plus its
Supplement) fram the general catalogue
of Ap and Am stars of Rensan (exclud-
ing those of peculiarity types Hg, Mn,
HgMn and MnHg, which show less or no
variability). Figure 2a (top) shows the
number of stars with determined
periods as a function of Right Ascen-
sion. The cross-hatched area concerns
the unambiguous periods, the hatched
area indicates uncertain or ambiguous
periods and white areas are for constant
or suspected long period stars, for
which no period could yet be found.
Figure 2b (bottom) shows the fraction of
Ap stars with known periods. From 250
bright stars, 32 per cent have a well-
defined period, and this percentage in-
creases to 42 per cent if uncertain or
ambiguous periods are included, as weil
as stars with very long period or no
variation. A smaller number of stars that
are observable in the southern winter
have been sufficiently studied to allow
the determination of the period; the con-
trast is, however, less striking in terms of
percentage than of total number, due
to the distribution of known Ap stars
acrass the sky.
In addition, the number of periods
known for cluster or association mem-
bers now reaches about 40. This allows
to see how the rotation period varies
with age, and it seems that it decreases
at the exact rate expected from conser-
vation of angular momentum during the
stellar evolution on the main sequence.
It is on the other hand noteworthy that
among those stars with well-determined
periods, by far the largest fraction are
CP 2 stars with enhanced Si or CP 4
stars. This is quite certainly a selection
effect, since these stars have periods
which very rarely are langer than 10
days, in contrast to the CP2's with en-
hanced Sr, Cr and/or Eu, among which
about 25 per cent have periods in ex-
cess of 10 days.
z 6
0 2 4 6 8 10
Period (days)
Figura 1: Histogram of the periods deter-
mined since 1983. Some periods are uncer-
tain, especially three of those that are smaller
than 0.5 days. Three more stars have periods
longer than 10 days.
Let us now discuss some interesting
aspects related to photometrie variabili-
ty. Mathys et al. (1986) present 12 B to
early F type stars which proved variable
(they were used as camparisan stars),
and five of them show strictly periodic
variations which are reminiscent of CP
stars. Their spectral type, however, is
not known to be peculiar, although all
have an MK c1assification (four of them
are in the Bright Star Catalogue). One of
these stars (HO 90994) is even a primary
standard for MK classification! It is very
interesting to note that three of these
2 4 6 6 10 12 14 16 16 20 22 24
Rlabt ••0000100 (h)
Figura 2 a (top): Histogram of the periods de-
termined up to now for bright non-HgMnAp
stars as a funet/on of right ascension. Cross-
hatehed area: well-defined periods. Hatched
area: uneertain or ambiguous periods. Blank
areas: stars for whieh no variation eould be
2b (bottom): Pereentage of bright Ap stars
with known period (compared to the total
number of known bright Ap stars) as a fune-
tion of right aseension. Same eode as in 2a.
several sueeessively excited pulsation
modes, one can derive important para-
meters, like mass and age, internal rota-
tion, depth dependeney of a magnetic
field, loeation of a eonveetive zone and
the radial dependency of temperature
and density. Thus, astroseismology
offers a unique way to look into stars
and to check the theory of stellar evolu-
tion. Further contributions at ESO from
our WG: Weiss and Schneider (1984)
observed the brightest member of this
group, a Cir, with the Walraven photo-
meter attached to the 90-cm Outch tele-
scope. For the first time, it was possible
to investigate amplitudes and phase
shifts of pulsation modes in a large
spectral range. During a compaign,
coordinated by Don Kurtz in South Afri-
ca (SAAO), Schneider and Weiss (1986)
observed HR 1217 and were able to
restriet possible mode identifications.
They also took part in a second interna-
tional observing campaign for HR 1217
in 1986 which resulted in a pulsation
frequency speetrum with the highest yet
known aecuraey (down to the noise limit
of 0.1 mmag), besides the sun (submit-
ted). Baade and Weiss computed syn-
thetic spectral line profiles for non-ra-
dially pulsating stars, taking different as-
pect and pulsation parameters into
aceount (1987). Simultaneous spectros-
copic and photometrie observations of
a Cir, obtained by Schneider and Weiss
(1989), yielded an upper limit for possi-
ble radial velocity variations of 100 m/s
in contradietion to 1 km/s elaimed by
other authors, although the star was
pulsating with an amplitude of 6 mmag
in Stroemgren v. A similar paper on
y Equ is presently being prepared for
publication by the same authors. Last
year, Schneider and Weiss participated
in an international observing eampaign,
• o
tem because it shows extreme photo-
metrie peculiarity. Surprisingly, its
period is as short as 0.476 days, and it is
probably the shortest, non-ambiguous
period among the Si stars.
Finally, two stars have at least two
periods: HD 25267 (t9 Eri) and HO
37151. They might be intermediate
cases between Ap stars and the mid-B
variables, one period being due to abun-
dance inhomogeneities and rotation, the
other period possibly due to non-radial
Figure 3: Lighteurve ot HO 187474 in Strömgren y. Open symbols reter to LTPV-group
observations, elosed symbols to older observations that have been shifted by 6.3 or 12.6
years. Different shaped symbols reter to the teleseope + photometrie system: ESO 50-em
(eirele), Danish 50-em (square), Boehum 60-em (diamond), all at La Silla, and eTio 1-m
,.- -0.96
3. CP2 Stars and Astro-
Ten years ago, Don Kurtz announced
the diseovery of light variations with a
period of about 12 minutes for the pecu-
liar CP2 star HD 101065, also known as
"Przybylski's Star". Weiss and Kreidl
(1980) observed this object at the same
time at ESO with Hß photometry and
presented the first evidence that pulsa-
tion is the relevant meehanism for its
photometric variability.
Observing stars of this class means to
detect light amplitudes which are fre-
quently of the order of less than one
millimagnitude. Therefore, even for
bright stars one needs teleseopes of the
1-metre class in order to reduce the
photon noise and the scintillation noise
due to our terrestrial atmosphere during
integration times of only 10 seconds
(periods range from about 4 minutes to
15 minutes!). By speetroseopie obser-
vations, in particular radial velocity mea-
surements, one ean see a star pulsating,
but obviously still larger telescopes are
required to aehieve the desired S/N ratio
from integration times of one minute at a
spectral resolution of R = 100,000.
For stars where one is able to observe
five stars have a rather early spectral
type, Le. B6V, and that all five stars
have periods smaller than 3 days, apart
from one (uncertain!) exception. This im-
plies that, although the periods and
Iightcurves of these stars resemble
those of c1assical CP stars, their proper-
ties are equally, and perhaps even more,
reminiseent of the so-ealled variable
mid-B type stars, discussed by Wael-
kens and Rufener (1985) as probably
being non-radial pulsators, like the 53
Per variables. The lack of any photomet-
rie peculiarity (indicated by Maitzens's
L\a or Geneva L\ (V1 - G) parameters)
would tend to eonfirm this analogy, as
weil as the fact that Waelkens and Ru-
fener (1985) eonsidered HO 74560 as
typieal "mid-B variable" while it is a
B4pMgSi star. The photometric varia-
tions of the latter star would undoubted-
Iy have been interpreted in the frame-
Work of the Oblique Rotator Model by an
Ap specialist!
Similarly, it can be questioned
whether the variability of Hg Mn Bp stars
should not be attributed to non-radial
pulsations rather than explained by the
Oblique Rotator Model.
Another interesting feature in the pro-
gress made sinee 1983 is the quantita-
tive deseription of the lighteurves. North
(1984) has shown that practically all
Iightcurves of CP stars can be satisfac-
torily represented by a sinusoid and its
first harmonic. This was confirmed by
Mathys and Manfroid (1985), who pre-
Sented a large, homogeneous sampie of
uvby lightcurves for 56 stars and listed
the fitted Fourier eoefficients. The fact
that no higher frequency is needed to
represent these lightcurves suggests
that there are only two important spots
(or, conversely, a unique ring) with
anomalous abundances on the surface
of Ap stars.
A few especially interesting stars
should still be mentioned. The long
Period star HO 187474 has been moni-
tored in Strämgren photometry, in the
framework of the long-term photometrie
programme of the Sterken group at
ESO. Figure 3 shows the visual Iight-
eurve of this star which is a famous
speetrum and magnetic variable with a
period of about 6.3 years. Data before
1983 (i. e. before the Sterken group) are
probably less homogeneous. A number
of other CP stars with suspected long
periods are regularly observed within
the same project.
The stars HO 98457 and HO 57946
have an extreme photometrie am-
Plitude, Le. about 0.2 mag in the near
ultraviolet. This is comparable to the
amplitudes previously found for HO
215441, GC 17353 (CpO-55 5216) and
for HO 187473. The SiMg star HD 60431
has been monitored in the Geneva sys-
similar to that for HR 1217. The results
for HD 203932 are being prepared for
4. Main Sequence Evolution:
Search for CP2 Stars in Open
The study of CP stars in open clusters
yields information and constraints on
the time dependence of the Ap
phenomena. During the discussions in
the first two meetings of the WG (Vienna
and Paris 1979) it became clear that one
has first to identify CP stars in open
clusters as completely as possible and
to enlarge the sampie of open clusters in
order to derive statistically sound re-
sults. Hitherto published spectroscopic
surveys were regarded as insufficient
concerning both requirements men-
The way out of the limitations of the
time consuming, subjective spectros-
copic surveys was provided by the use
of photometric indices sampling the Ap
typical broad band flux depression
around 5200 Ä. The first observing runs
at the end of 1979 at the ESO 50-cm
telescope by Maitzen and Hensberge
(1981) used the L1a-index (Maitzen,
1976) to pick out CP2 or CP4 stars in
NGC 2516 and NGC 1662. Maitzen and
Vogt (1983) demonstrated with a large
sampie of field Ap stars that L1a has the
same detection capability for magnetic
Ap stars as the high qualitiy classifica-
tion dispersion spectroscopy of Bidel-
man and MacConnell (1973).
Meanwhile 13 papers (for references
see Maitzen et al., 1988) have been pub-
Iished on 33 open clusters with slightly
more than 1,000 stars surveyed, most of
them after 1983. Observations for
another 32 clusters have been obtained
and are being evaluated.
Our limiting magnitude has been set
by telescopes in the range 50 to
100 cm, and is about 11-12 mag, re-
spectively. The number of clusters
where a sizable number of objects (i. e.
more than 10) can be surveyed in the
domain B5-A5 with these limits, is
roughly 80. Thus, with our technique
and telescopes available we are ap-
proaching the end of our programme.
The relatively large number of clusters
in our survey is justified by a preliminary
evaluation of the results which do not,
contrary to Abt (1979), exhibit a simple,
clear-cut picture, i. e. dependence with
time. Other parameters are probably
playing a role in determining how many
CP2 stars are present in a cluster and
how strong their peculiarities will de-
velop. A careful statistical examination
of the data (which are much more
homogeneous than spectroscopic re-
sults) to be carried out in our WG during
this year will try to find out whether there
is now enough evidence for statistically
sound conclusions. If not, fainter clus-
ters may be included using larger tele-
scopes and CCD techniques. Our actual
sampie should comprise of the order of
100 peculiar stars, taking into account
the frequency of field CP2 stars (about 5
per cent).
The outstanding role of ESO instru-
ments is underlined by the fact that
more than half of all observations have
been made using either the ESO 1-m,
ESO 50-cm or Bochum 61-cm tele-
scopes. This percentage rises to 87, if
one considers only the southern hemi-
sphere observations. So far the follow-
ing participants have contributed to the
project: H. Schneider (Göttingen), H. M.
Maitzen (Vienna), K. Pavlovski (Zagreb),
H. Jenkner (Baltimore-ESA), F.A.
Catalano (Catania), H. Hensberge
(Brussels), W. W. Weiss (Vienna), T.
Kreidl (Flagstaff), G. Deridder (Brussels),
M. Floquet (Paris), H. J. Wood (La Sere-
na-CTIO) and C. Tanzer (Vienna).
5. Magnetic Fields
One of the most remarkable charac-
teristics of the Ap stars of the Sr-Cr-Eu,
Si, He weak and He rich types is that
they possess a magnetic field with a
large-scale organization, a unique prop-
erty among nondegenerate stars. The
presence of an organized magnetic field
is thus intimately related to the appear-
ance of other peculiarities, in which it
most probably plays a key role. More-
over, the magnetic field is likely to have
a significant influence on the observed
stellar properties, so that its knowledge
is an essential part of the proper under-
standing of the physics of Ap stars. (For
instance, the magnetic field must be
taken into account when deriving ele-
mental abundances).
Since 1984, a Zeeman analyzer is
available as standard option on the
CASPEC at the 3.6-m telescope, pro-
viding European astronomers with the
opportunity of studying magnetic fields
of Ap stars from circular polarization
measurements in their spectral lines. An
observing programme aiming at the de-
termination of the spatially unresolved
structure of the magnetic field of a
number of stars has been initiated in our
WG by Mathys, taking advantage of the
unprecedented possibility offered by the
Zeeman analyzer of the CASPEC to re-
cord simultaneously the profiles in cir-
cular polarization of a large number of
spectrallines, with high S/N ratio. This is
a long-term project, since spectra suit-
ably distributed over a stellar rotation
cycle are needed in order to untangle
the geometry of the field. The acquisi-
tion of the data is now mostly com-
pleted, for 9 stars a good phase cov-
erage could be achieved; more partial
data have been obtained for 25 addi-
tional stars. Their analysis is in progress.
Spectra recorded through a Zeeman
analyzer essentially contain information
about the component of the magnetic
field vector along the line of sight. By
observing on very high resolution spec-
tra (recorded without apolarimeter) the
subtle broadening induced in the spec-
tral lines by the magnetic field, it is pos-
sible to study its modulus. This informa-
tion complements that obtained through
spectropolarimetry and thus permits to
set more constraints on the stellar
magnetic field. The interpretation of the
magnetic broadening, however, is not
straightforward. A programme is being
carried out by Didelon in order to deter-
mine whether the Robinson (1980) tech-
nique, which has been successfully ap-
plied to magnetic field determination in
late-type stars, can be employed to
measure magnetic fields in Ap stars.
The principle of the technique is to com-
pare the profiles of two lines, one which
is sensitive to the magnetic field and the
other which is not, in order to derive the
magnetic field from the difference in
their widths. Didelon's observations are
being made with the CES at the CAT.
The magnetic field determined with
the Robinson technique from unsplit
lines is compared with that obtained
from lines which are fully split into their
Zeeman components by the magnetic
field. Not that the latter effect can be
observed only in a few stars where the
relevant parameters (magnetic field, ro-
tational velocity, angle between the ro-
tation axis and the line of sight) have
quite favourable values. One can advan-
tageously use these stars to calibrate
the Robinson technique, which can later
be applied to determine magnetic fields
in less favourable, more widespread
cases. A progress report on this project
was published recently in the
Messenger No. 49 (Didelon, 1987).
6. Spectroscopic Studies
The main characteristics defining the
various types of CP stars is the anomal-
ous strength of some of their spectral
Iines. These spectral peculiarities refleet
the overabundance or underabundance
of some chemical elements in the stellar
atmosphere. Abundance determinations
are thus an essential part of the study of
Ap stars, all the more since they carl
strongly constrain the theoretical mod-
els which are developed to explain the
origin of these objects.
Spectroscopic investigations in the
region covered by the International VI-
traviolet Explorer (tUE) have been per-
formed for a number of CP stars. Most
of the concerned spectra have their re-
sonance lines in this range. However,
line identification is generally hampered
by the high density of competing tran-
Fuhrmann devoted special interest to
the well-known CP2 star HR 465, while
others like HR 7775, HR 4072, x Cnc,
a2CVn, HO 101065, as weil as their
"normal" congeners j[ Cet, v Cap and
Vega were preferably used for compari-
son purposes. The spectra show a huge
number of absorption lines, most of
which belong to Fell and Crl!. Other
elements - like Mg 11 or Si 11 - have only
few, but strong transitions. In the spec-
tra of the CP2 star a2CVn a remarkable
number of ions with Z < 20 is not weil
pronounced. This is most obvious in the
case of neutral and ionized carbon.
As far as the Rare Earth elements are
concerned there are identifications for
some second spectra (e.g. HolI, Gd 11,
. ..) in the tracings of HO 101065 -
Przybylski'S star. The spectra of the
somewhat horter HR 465, however,
show only marginal contributions from
this group of elements. IUE spectra are
Weil suited to show the definite pres-
ence of heavy elements like platinum
and mercury, as weil as the overabun-
dances proposed from optical spectra.
Additionally, absorption lines of Bill are
observable. There is also strong evi-
dence for transitions due to gold (Au 11).
In the visible, Gerbaldi and Faraggia-
~ a have been extending the investiga-
~ I o n of abundances to elements not eas-
~ observable with photographic plates.
hey have derived intriguing results
from observations of the neutral lithium
resonance line at A. 6707 now easily ac-
cessible with modern detectors. The
first observations made by them at ESO
with CES and Reticon of the Li I 6707
line in cool Ap stars, raised several
problems since the feature at this
wavelength shows an asymmetric pro-
file different from star to star and an
intensity which is not related to the at-
mospheric parameters of the stars. Sur-
prisingly, the only rough relation de-
tected was that between the equivalent
width of A. 6707 and the number of
other, mainly unidentified lines present
in this spectral range (Gerbaldi,
Faraggiana, 1986). Subsequent obser-
vations at the Observatoire de Haute-
Provence, complementary to those per-
formed at ESO, indicate that a line of
another element, so far unidentified, is
present at a wavelength very close to
the Li line.
Spectroscopic studies of HgMn (CP3)
stars were carried out by Schneider
(1986), who showed that more than 60
per cent of these stars are binaries. Us-
ing new observations he raises this val-
ue to more than 70 per cent. This fact
puts the CP3 group close to the CP1
(Am) stars, in contrast to the CP2 stars
with a binary frequency of only about 30
per cent. The CP3 stars show a con-
centration towards circular orbits for
short periods and a lack of periods less
than two days. Stars with periods less
than 10 days tend to synchronized rota-
tion. They are all slow rotators with v·sini
of about 30 km/so
Abundance analysis of some CP3
stars were carried out by Ansari who
carefully studied the effect of rotational
broadening of spectral lines on the de-
rived abundance values.
Abt, H.A.: 1979, Astrophys. J. 230, 485.
Baade, D., Weiss, W. W.: 1987, Astron. Astro-
phys. Suppl. 67, 147.
Bidelman, W.P., MacConnell, D.J.: 1973, As-
tron. J. 78, 687.
Catalano, F.A., Renson, P.: 1984, Astron.
Astrophys. Suppl. Sero 55, 371.
Catalano, F.A., Renson, P.: 1988, Astron.
Astrophys. Suppl. Sero 72, 1.
Didelon, P.: 1987, The Messenger49, 5.
Gerbaldi, M., Faraggiana, R.: 1986, Astron.
Maitzen, H. M.: 1976, Astron. Astrophys. 51,
Maitzen, H.M.: 1981, Astron. Astrophys. 96,
Maitzen, H. M., Schneider, H., Weiss, W. W.:
1988, Astron. Astrophys. Suppl. Sero 75,
Maitzen, H. M., Vogt, N.: 1983, Astron. Astro-
phys. 123, 48.
Manfroid, J., Mathys, G.: 1985, Astron. Astro-
phys. Suppl. Sero 59, 429.
Manfroid, J., Mathys, G.: 1986, Astron. Astro-
phys. Suppl. Sero 64, 9.
Mathys, G., Manfroid, J.: 1986, Astron. Astro-
phys. Suppl. Sero 60, 17.
Mathys, G., Manfroid, J., Renson, P.: 1986,
Astron. Astrophys. Suppl. Sero 63, 403.
North, P.: 1984, Astron. Astrophys. Suppl.
Sero 55, 259.
North, P.: 1987, Astron. Astrophys. Suppl.
Sero 69, 371.
Robinson, R.D.: 1980, Astrophys. J. 239,
Schneider, H.: 1981 , Astron. Astrophys.
Suppl. Sero 44, 137.
Schneider, H.: 1986, "Upper Main Sequence
Stars with Anomalous Abundances", IAU-
Coll. No. 90, eds. C. R. Cowley et al. ,
Schneider, H.: 1987, Hvar Obs. Bull. 11,29.
Schneider, H., Weiss, W. W.: 1986, Mon. Not.
Roy. Astr. SOG. 215, 77.
Schneider, H., Weiss, W.W.: 1989, Astron.
Astrophys., in press.
Weiss, W. W., Kreidl, T. J.: 1980, Astron. As-
trophys. 81, 59.
Weiss, W. W., Schneider, H.: 1984, Astron.
Astrophys. 135, 148.
Waelkens, Ch., Rufener, F.: 1985, Astron.
Astrophys. 152, 6.
AK Scorpii: A NewPre-Main-Sequence Spectroscopic Binary
J. ANDERSEN, Copenhagen University Observatory, Denmark, and Center tor Astrophysics,
Cambridge, Mass., USA
H. LlNDGREN, ESO
M. L. HAZEN, Center tor Astrophysics, Cambridge, Mass., USA
M. MAYOR, Observatoire de Geneve, Switzerland
While we know that perhaps - 25 %
of normal main-sequence stars are
spectroscopic binaries, very few of their
Progenitor systems have yet been de-
tected among (Iow-mass) pre-main-se-
quence stars: Mathieu's (1988) review at
the recent lAU General Assembly lists
only 11, and for only three had orbits
been published at that time. This
meagre yield must be due mainly to
selection effects mitigating against
the discovery of pre-main-sequence
binaries. These stars are intrinsically
faint and generally found in highly ob-
scured regions, so systematic and
accurate radial-velocity observations
were impossible until the advent of effi-
cient cross-correlation techniques. Oue
to their importance for the understand-
ing of star formation processes in gener-
al, pre-main-sequence binaries are now
being searched for very actively, and the
sampie will no doubt increase sharply
over the next few years. We would like
150 ...---r--"T1----,1--r-.-"T1----,1--r-.-"T1----,1--r---
-..,....- 1-.....
The Environment of AK Seo
Although, speetroseopieally, AK Seo
itself looks rather ordinary, it appears to
live in quite a Iively plaee, judging by its
irregular variability, which seems too
large ("" 1 mag) to be reasonably ex-
plained by star spot aetivity. The UBVRI
photometry of Kilkenny et al. (1985)
shows that the star gets redder as it
gets fainter, with a ratio of total to selee'
tive absorption R = A
E(B-V) "" 4.6. This
is mueh higher than in ordinary interstel-
lar matter (R = 3), as often oeeurs in
eireumstellar dust and is usually eonsid-
ered due to a relatively large grain size.
There is additional evidenee for eir-
eumstellar dust from infrared photome-
try of AK Seo: Figure 2 shows the ener-
gy distribution derived from the
UBVRIJHKL photometry by Kilkenny et
al. (1985) and the 12, 25, and 60 f-tm
IRAS fluxes. The fluxes have been
eorreeted for an average amount of ex'
tinetion. Three blaek-body eurves have
been fit to the data. In addition to the
flux from AK Seo itself, they show the
F5 main-sequenee stars. Popper's
(1980) review indieates that reasonable
ZAMS masses and radii are 1.3 MG and
; the orbital inelination is then
i"" 69°. If the orbital and axial rotation
periods are equal, we expeet to mea-
sure rotational veloeities of vsini ""
4 km S-1. However, due to tidal effeets,
eonveetive stars in eeeentrie systems
are expeeted to rotate faster than this,
about 2.5 times faster for the orbital
parameters of AK Seo. This revises our
predietion to vsini "" 11 km S-1 for both
What we aetually measure from the
width of the GORAVEL eross-eorrelation
profiles is vsini = 19 ± 1 km S-1, so the
stars either spin faster or have larger
radii than first assumed. Gonsideration
of the time-seales for synehronization
suggests that the former is unlikely,
espeeially sinee the stars were probably
even larger and easier to synehronize
when they were younger. Henee, the
stars in AK Seo are probably weil above
the ZAMS.
Published models of pre-main-se-
quenee evolution are relatively old and
still largely untested, so we searehed
instead in Popper's (1980) review for a
real, suitably evolved binary as a "role
model" for AK Seo. The F5-type system
RZ Gha (1.5 MG, 2.2 R
), an old friend
of ours from early days on La Silla,
seems to fit the bill (and the rotation) for
an inelination of i"" 63°. In this model,
the stars still narrowly fail to eelipse.
Assuming an effeetive temperature of
6,500 K, AK Seo is about 200 pe distant,
eonsistent with membership in the Seo'
Gen assoeiation.
1.0 0.9 0.8 0.7
A Model of the Binary
If we eannot aetually determine the
masses and radii of AK Seo, what might
they reasonably be? Speetroseopieally,
the eomponents of AK Seo look just like
1.5-m teleseope on La Silla and deter-
mined its speetroseopie orbit as shown
in Figure 1. The orbital eeeentrieity is e =
0.47, the period is 13.6093 days, both
minimum masses (msin
/) are
1.06 ± 0.01 solar masses, and the orbi-
tal semi-axis major is about 31 solar
radii. The origin of the phases plotted in
Figure 1 is where a primary eelipse
would oeeur; the phases of a possible
seeondary eelipse and of periastron
passage are also indieated.
Unless we ean determine the inelina-
tion of the orbit, we eannot find the
absolute masses of its eomponents,
and unless the system eelipses, we
eannot direetly determine their radii
either. Now that we know when to look
for eelipses in AK Seo, do they in fact
oeeur? We searehed for the answer in a
nearby "gold mine", the vast Harvard
plate eolleetion going back about a een-
tury. And yes indeed, it holds more than
2,000 plates on whieh AK Seo ean be
measured! The period 1910-12 had a
partieularly dense eoverage of plates
(83), so we estimated the brightness of
AK Seo on these plates and plotted the
resulting magnitudes against speetro-
seopie phase. Alas - no eorrelation
whatever! So, while shallow eelipses
might still be diseovered by eareful
photoeleetrie photometry. most of the
-1-magnitude variations we do see
must have some other origin, probably
in nearby dust elouds.
/ \
/ \
, , ------.---
,/ /
..... '" .,
100 r-
i n
-150 L-_.J.- 1_....I•.L-t_LJI!L..-_.J.-I_....II__ _ ___L_---'
0.0 0.1 0.2 0.3 0.4 0.5 0.6
Figure 1: Speetroseopie orbits tor AK Seo. Phases are eounted trom a hypothetiea/ primary
eelipse; the phases ot seeondary eelipse (/I) and periastron (P) are indieated. Cross: eoude
speetrum; dots: CORAVEL observations; eire/es: sing/e-/ined CORAVEL observations. not
used in the solution.
to share with readers some of the fun we
have had with one of these new sys-
tems, AK Seo (full details will appear in
Astron. Astrophys.).
AK Seorpii: Classieal or "Naked"
T Tauri Star?
AK Seorpii is an 8th magnitude mid-F
type variable star projeeted on the
southern outskirts of the Seo-Gen
assoeiation. Herbig and Rao (1972)
found it to have a strong Li I line at
6707 Aas weil as strong Ha emission
with a eentral reversal, showing it to be
very young. It shares its large-amplitude
light variations and infrared exeess with
the "elassieal" T Tauri stars, but the
absence of emission lines other than Ha
makes its optieal speetrum more similar
to those of the "naked" T Tauri stars
(Walter et al. 1988), which may owe their
lack of eonspieuous eireumstellar mat-
ter to the presenee of a binary eom-
A eoude speetrum of AK Seo taken
with the ESO 1.5-m teleseope in May
1986 showed it to be a double-lined
speetroseopie binary with approximate-
Iy equal eomponents. As AK Seo had
onee been eonsidered to be an eelipsing
binary, we hoped that it would be possi-
ble to determine the masses, radii,
luminosities, and eomposition of its
eomponents. Gomparing them with stel-
lar evolution models, age, helium abun-
dance, and other interesting parameters
might then be determined.
km/s 0 I-
Orbital Parameters
We therefore observed AK Seo in
1986-87 with GORAVEL on the Danish
0.80 '----,_ ,...--'--7:
Figure 3: Cross-eorrelation profiles for two
CORAVEL observations (phases 23 and
48), separated by only 1.2 orbital eyeles.
. ...
1.00 &:tp..... ._4iflt-.
, ""'.
r •....W"'·
0.95 .. .t
0.90 f--t--t--t--t--t--+--+--+--+-+--+--+---J
Wavelength (JJ-)
Figure 2: Blaekbody fit to the optieal-infrared energy distribution of AK Seo. The three eurves
eorrespond to temperatures of 6,500, 1,600, and 160 K.
10-l2 L..l_..l-..lL....L.Ju---W...__-----"l-----'_-'---L--..L.C.JL.L.L..L \..:L_....L---L-----l---l-l
presence of two dust components at
approximate temperatures of 1,600 K
and 160 K. The total energy emitted in
these components amount to about half
of that received from AK Sco, roughly
consistent with the light loss at inter-
mediate brightness levels.
Size of the Dust Clouds
We can obtain an order-of-magnitude
estimate of the sizes of the two dust
cloUds by requiring that the energy emit-
ted by a grain be equal to that it absorbs
from the star. The resulting distance
from AK Sco is then 25-50 solar radii
for the hot dust, and about 10 AU - the
orbital radius of Saturn - for the cool
component. Thus, the hot dust cloud
has essentially the same size as the
binary orbit and presumably consists of
material lett over from its formation,
while the cool dust c10ud has the size of
a typical planetary system (assuming
that the one in which we live is typical!).
The CORAVEL data contain one strik-
piece of information supporting this
PICture: The luminosity ratio between
the two components, as measured by
the equivalent widths of their cross-cor-
relation dips, changes up or down by up
to a factor two from one orbital cycle to
the next (Fig. 3). Again, star spots are
unlikely to cause such large variations;
even if they did, they should significantly
distort the profiles, which is not seen.
Moreover, the variations are largest
when the stars are farthest apart.
The simplest interpretation appears to
be that we are seeing the effect of in-
homogeneities in the hot dust c10ud not
only in front of, but probably also be-
tween the stars. This opens an exeiting
possibility for mapping out the dust
clouds near the system, in the following
way: At any given time, we know the
preeise position of both stars in their
orbit. Photometry teils us the total ob-
seuration in front of the two stars, and a
simultaneous CORAVEL or other spec-
troseopie observation ean give us the
ratio of the light losses for each compo-
nent. We can therefore compute, point
by point, the amount of dust in front of
each star as it goes through an orbital
eycle, and draw a map of the dust
elouds and eventually of their motions.
Perhaps one could even measure the
velocities of eircumstellar Iines at high
Thus, AK Sco no doubt has lots more
fun in store for the observers: There are
eclipses to look for and dust c10uds to
map. Spectroseopically, one can inves-
tigate the Li and other element abun-
dances, indicators of ehromospherie
activity, and signatures of possible
magnetic fields: AK Sco is within reach
of the CES. And perhaps, one day, it will
be possible to resolve the system inter-
ferometrically (with the VLT?), although
the angular separation is only of the
order of 1 mas; its absolute masses and
distance could then be determined even
if eelipses should not occur.
Herbig, G.H., Raa, N.K.: 1972, Astrophys. J.
Kilkenny, D., et al.: 1985, South Afr. Astron.
Obs. Cire. No. 9, 55.
Mathieu, R.D.: 1988, Highlights of As-
tronomy, Val. 8, in press.
Popper, D. M.: 1980, Ann. Rev. Astron. Astro-
phys. 18, 115.
Walter, F.M., et al.: 1988, Astron. J., 96, 297.
SupersheIls and Galactic Fountains
, Astronomical Council, Moscow, USSR
In the gaseous disk of our Galaxy as
Weil as in other galaxies, HI structures
(shells, bUbbles, holes, ete.) on seales of
0.1-1 kpc are recognized to be com-
-1 Visiting astronomer, ESO _ Garehing.
mon features; see e. g. the comprehen-
sive review by Tenorio-Tagle and
Bodenheimer (1988). The larger ones
are usually named with the prefix "su-
per". The estimated energies which are
required to produee sueh large objeets
are high - up to some 10
erg. These
energetic events must exert a significant
influenee upon the gaseous galactic
disk and eorona.
In our Galaxy, the disk and the corona
are believed to be evolutionarily coupled
in a recycling process, although there is
no eommon opinion about details of the
Figura 1: Evolution of a supershell tE =2 X 10
erg S-I) in a gaussian galaetie disk, with no = 1 em-
, H = 140 pe. Density eontours are shown
with logarithmie spaeing of 0.5 dex at times 0.2, 0.8, 1.5,2.1,2.5,3.0 Myrs. The maximum veloeity is 2,200 km S-I.
process. An evaluation of the density n
and temperature T in the corona and
elucidation of the distributions n(z) and
T(z) in the disk remains achallenging
subject for observers and theoreticians.
It would appear that it is possible to
arrive at a better understanding of the
topic by interpreting the high energy
processes in the disk and corona in a
consistent way. The construction of
theoretical models of large scale HI
structures in galaxies may be consid-
ered as a step in this direction.
Several mechanisms have been pro-
posed to explain large scale HI struc-
tures in galaxies. The hypothesis of
supernova explosions appears to be
most consistent with the idea of heating
and supporting of the corona by galactic
sources (Cox and Smith 1974). How-
ever, a single supernova event seems to
be unable to break out of the HI disk
unless the explosion is presumed to
take place substantially off the plane.
This means that single explosions can
hardly be considered as a source of hot
gas for replenishing the corona. Multiple
supernova outbursts are more probable,
since OB stars are weil known to be
born in groups, which in the solar vicini-
ty on the average contain some tens of
these stars (Ambarzumian 1947, Blaauw
In order to investigate the evolution of
(super)shells which originate as a pro-
duct of interaction of supernovae with
interstellar matter, we have calculated
2D hydrodynamical numerical models.
Some important values that are not weil
known from observations, like energy
input rate E, parameters of gas distribu-
tion in z-direction, etc., were considered
as model parameters. A typical value of
Efor an OB association in the Galaxy
seems to be of the order of 10
erg S-1.
According to the models, a break-out
from the gaseous galactic disk occurs in
the course of the evolution of an OB
association in the plane of the disk with
gaussian density distribution and densi-
ty scale height H = 140 pc for E =
2 x 10
erg S·1 in - 7 x 10
years. For
the more energetic case (E = 2 x 10
erg S·1) this event occurs after
- 2.5 x 10
years. In Figure 1, the evolu-
tion of the latter model is illustrated.
The first phase of this evolution is
rather weil understood in terms of the
theory of interstellar bubbles driven by
stellar wind (Weaver et al. 1977). Super-
nova explosions are interpreted in our
models as a continuous and powerlul
stellar wind. The freely expanding wind
soon stops and the velocity of expan-
sion decreases from - 2,000 km S-1 to
-100 km S·1. The stopped wind then
acts as a giant, hot isobaric gaseous
piston driving a massive cool shell. A
dense, relatively cool HIsheIl contains
swept interstellar gas. This gas is shock-
ed (the shell velocity is much greater
than the sound speed in an ambient
interstellar medium) and heated up to
K. Then it cools down to T == 10
or less, due to radiative losses of inter-
nal energy. The top of the shell becomes
unstable (Rayleigh-Taylor instability)
when it reaches - 1.5 Hand then breaks
into fragments. The fragments can be as
massive as -10
solar masses and
move with -100 km S-1. For the lesS
energetic case (typical for the Galaxy),
the masses and velocities of these frag-
ments are 2-3 times lower. The number
of fragments is of the order of ten, and
the total mass never exceeds 0.1 times
the mass of the shell. Calculations with
an exponential atmosphere reveal a
smaller tendency to fragmentate. As
said above, however, it is not yet pos-
sible to make an appropriate choice of
the galactic atmosphere. Moreover, the
structure of the disk in the z direction
appears to be dependent on blow-out
events and the following dynamical
evolution of the expelled gas. Corbelli
and Salpeter (1988) have for instance
shown that relatively weak current blow-
out activity can be eftective in compres-
sing an outer HI disk and, possibly, in
giving it a sharp edge.
The evolution of the expelled gas may
be described in terms of a "galactic
model, as proposed by Shap-
Iro and Field (1976). Due to thermal in-
stability, the rising hot gas converts into
HI clouds falling onto the disko These
clouds resemble some of the so-called
"high velocity" clouds (HVC) discovered
by Dutch astronomers in the middle of
the 60's; see the recent survey paper of
Hulbosh and Wakker (1988). Such mod-
els of the galactic HVC - the first one
Was elaborated by Bregman (1980) -
seem to be attractive. The main problem
of the models is poor knowledge of n
and T in the galactic corona and conse-
quently a badly known cooling time
Scale. For typical T = 2 x 10
K and n =
2 x 10-
, say, the (radiative) cool-
ing time exceeds 10
years, which
appears to be Ionger than the time scale
of heating by supernovae.
Without discussing in detail the re-
sults obtained in our model calculations
(see Igumentshchev et al. 1988, 1989;
very recent results are in preparation for
publication), we note that cool frag-
ments may be interpreted as seed
clouds for HVC. The masses of ob-
served HVC, when estimated under the
assumption that they are situated at
mean galactic distances, are 10
solar masses. The fragments seem to be
smaller, but it is possible for them to
grow. They are massive and large
enough (R 20-50 pc) to persist
against evaporation. Moreover, if rising
with the given velocities up to
0.7-2 kpc over the plane, they could
act as a sink of thermal energy of hot
coronal gas. A simple estimate, accord-
ing to the theory of McKee and Cowie
(1977), shows that the largest fragments
are resistant to evaporation (radii ex-
ceeding the critical value) and should
Note also that the proposed mechan-
ism of sequential supernova explosions
cannot explain the extremely large (R
1 kpc) supersheIls unless one assumes
very low densities n 0.01 cm-
. Even
in that case one must suppose enor-
mously rich OB associations in order to
explain the large kinetic energies. On
the other hand, the most common
supersheIls with radii -0.2-0.5 pc are
explained quite weil by this model.
I am grateful to the staft of ESO for
help and friendly assistance during my
stay in Garching.
Ambarzumian, V.A.: 1947, Stellar Evolution
and Astrophysics., Armenian Aead. Sei.
Blaauw, A.: 1964, Ann. Rev. Astron. Astro-
phys. 2, 213.
Bregman, J.N.: 1980, Astrophys. J. 236, 577.
Corbelli, E., Salpeter, E. E.: 1988, Astrophys.
Cox, D. P., Smith, B. W.: 1974, Astrophys. J.
(Letters) 189, L105.
Hulbosh, A.N.M., Wakker, B.P.: 1988, As-
tron. and Astrophys. Suppl. 75, 191.
Igumentshehev, I. V., Tutukov, A. V., Shustov,
B. M.: 1988, Astron. Zirkular N 1532, 21.
Igumentshchev, I. V., Tutukov, A. V., Shustov,
B. M.: 1989, Astron. Zh. (in press).
McKee, C. F., Cowie L. L.: 1977, Astrophys. J.
Shapiro, P. R., Field, G. B.: 1976, Astrophys.
Tenorio-Tagle, G., Bodenheimer, P.: 1988,
Ann. Rev. Astron. Astrophys. 26, 145.
Weaver, R., MeCray, R., Castor, J., Shapiro,
P., Moore, R.: 1977, Astrophys. J. 218,
Searching for Light Echoes from Circumstellar Dust Shells
Around SN 1987A
F. PARESCE, Space Telescope Science Institute, Baltimore, Maryland, USA,
andAstrophysics Division, Space Science Department of ESA
t {he Space Telescope Science Insti-
e coronograph (described in the
2 essenger, 47, p.43) mounted on the
.2-m telescope at La Silla was used by
myself and my colleague Chris Burrows
for light echos in the immediate
of ,the SN 1987A in the LMC.
IS technlque allows us to probe cir-
cumstellar regions from about 2 to 20
in the vicinity of very bright
Jects such as the SN 1987A for faint
otherwise completely lost in
elr glare. At the time of the observation
here, carried out on December
E h1987, the SN brightness was V = 6.2.
l' c oes at these angular distances and
of observation correspond to linear
IS ances from the SN roughly in the
of 1-15 parsecs in front of the
the observer's line of sight. At
e dlstances, any echo would most
probably be emitted from matter ex-
pelled from the SN itself at some earlier
epoch when it was a hot main-sequence
giant and/or a red giant star in its way
towards eventually becoming aSN.
Theory predicts and many observations
confirm that shells of swept-up stellar
material including dust formed in the
outer layers of an expanding photo-
sphere or wind would be expected to
linger in the general vicinity of the SN
and be observable by the delayed scat-
tering of the SN light pulse. The echo
would manifest itself in the form of a
luminous ring of several arcseconds
radius centred roughlyon the SN itself.
The ring might, in practice, be incom-
plete if the shell structure were not
homogeneous. Rings of this type 10-
cated much further away have been ob-
served around SN 1987A (see the
Messenger, No. 52, p. 13) but are
thought to be due to scattering from
sheets of interstellar dust Iying between
us and the SN and not related to the SN
The accompanying composite image
corresponding to a 10-minute exposure
taken with the ST Scl coronograph
through a standard B filter i1lustrates
graphically the result we obtained. The
full field of view shown in this image is
22 by 36 arcseconds with North up and
East to the left. The occulted centre of
the SN is located at the position of the
cross at the centre of the image where
the occulting wedge running EW is 1.9
arcseconds thick. Each pixel in the
512 x 320 pixel frame corresponds to
an area of 5 . 10-
while 1"
corresponds to a linear scale of - 0.83
light-years at the source. The left or
eastern half of the image corresponds to
the eastern half of the cleaned, bias-
subtracted, f1at fielded image of the SN
and its surroundings. In this portion, the
seeing broadened image of the bright
central source extending to approxi-
mately 5 arcsec from the centre is very
prominent. The seeing during this expo-
sure was 0':8 approximately, as evi-
denced by the size of the stars in the
field and the narrow diffraction spike in
the SE.
The right or western half of the com-
posite image represents instead the
western half of the image obtained from
the first by careful subtraction of the
purely stellar profile. We found that a
comparison of the observed profile with
the theoretically expected one from at-
mospheric turbulence theory provides
an effective means of assessing the re-
sidual circumstellar emission down to
=2 arcseconds from the star. As can be
seen from even a superficial inspection
of the right half of this image, no con-
spicuous features appear anywhere in
the field beyond the dozen or so stars
and a few artifacts associated with the
detector. All the stars observed were
known to be there before the eruption
including star 2 of the Sk-69 202
plex located at 2':8 and 320
p. a. in the
image shown here. This is the first sight-
ing of this star after the SN 1987A erup-
tion and its magnitude (B = 15.1) con-
firms the pre-eruption measurements.
The rather surprising result we obtain is
that excluding the thin occulted strip
running EW through the image, the
presence of significant amounts of cir-
cumstellar material in the region probed
by the light pulse at the time of observa-
tion can be excluded to a reasonable
degree of confidence. Specifically, the
absence of detectable rings around the
SN to brightness levels greater than ap-
proximately 22.5 blue magnitudes arc-
implies upper limits to shell col-
umn densities of hydrogen ranging from
about 10
atoms cm-
in the innermost
observable regions to about 10
in the outer ones. If the shell is
broken up into smaller and isolated
components, these limits would have to
be revised upwards, of course, since
irregular diffuse features are harder to
Since the echo emission pattern
around the SN is very time dependent
owing to the rapid expansion of the
echo ellipsoid as it sweeps out into the
surrounding circumstellar medium, it will
be interesting to see whether this result
will continue to hold true. If it does, it
implies that this particular SN is located
in a cavity of suprisingly low gas density.
How this can come about is now a sub-
ject of some speculation. Continued
monitoring of the SN as it fades and the
delay ellipsoid sweeps into areas unob-
servable in the Oecember 21, 1987 time
period should provide exciting new in-
formation on the past evolution of this
amazing object.
Note added in proof
On February 6, H. E. Bond, N.
Panagia, R. Gilmozzi, and M. Meakes,
Space Telescope Science Institute, re-
port (lAU Circular 4733) that direct CCO
images of SN 1987A, obtained in B, V,
and R at Cerro Tololo on January 24,
1989, appear to show a new light-echo
feature. In addition to the two previously
known echo rings, which now lie at radii
of approximately 46" and 78" from the
supernova, a new, nearly circular feature
is seen at a radius of only 9':5 (projected
radius 2.4 pc). Its surface brightness is
comparable to that of the two outer
rings, but is highest in a rim on tM
eastern side of SN 1987A. This feature
would imply the existence of dust Iying
at a distance of 5.9 pc from the super-
nova (assuming that it is illuminated by
light emitted at the optical maximum in
May 1987).
The Geneva Photometrie Monitoring of SN 1987A
Geneva Observatory, Switzerland
' " ~ ; ' ' '
e , ~ , , _
.., ----,
Burnet, M., Rufener, F.: 1979, Astron. Astro-
phys. 74, 54.
Nomoto, K., Shigeyama, T., Kumagai, S.,
Itoh, M., Nishimura, J., Hashimoto, M.,
Saio, H., Kato, M.: 1988, to appear in
Physics of Neutron Stars and Black Holes,
ed. Y. Tanaka, Universal Academy Press,
Tokyo, p. 441.
Woosley, S.E., Pinto, P.A., Hartmann, D.:
1989, Astrophys. J., submitted.
SN 1987A Geneva Photometry
SN 1987A Geneva Photometry
12 LLJ-LL.J.-JL.L--LJL.L-LJ---L-l-l-L..L..L-L.L.LJ-LL-L.L.L--LJL.L-LJ-L..L..L-L..L..L-LJ
47200 47250 47300 47350 47400 47450 47500 47550 47600
Figure 2: Enlargement of the last portion of the Vcurve with magnitudes corrected for the two
companion stars. The dashed line is the same linear fit as in Figure 1. The fullline is the fit of a
parabola calculated for the values within the range HJO 2447270 to 2447450.
46800 46900 47000 47100 47200 47300 47400 47500 47600
Figure 1: The Vmagnitude curve, uncorrected for the contributions of the two faint companion
stars of Sk -69°202. The dashed curve is the linear fit in the range HJO 2446980 to 2447285.
Arnett, W. D.: 1988, Theoretical models of
supernova 1987A, in Supernova 1987A in
the LMC, eds. M. Kafatos and A. Michalit-
sianos, Cambridge University Press, Cam-
bridge, p. 301.
Burki, G., Cramer, N., Burnet, M., Rufener, F.,
Pernier, B., Richard, C.: 1989, Astron. As-
trophys. Letter, in press.
A systematie photometrie monitoring
of SN 1987A has been earried out from
the Swiss station at the ESO La Silla
Observatory (Chile), by using a 70-em
teleseope equipped with the P7 photo-
meter (Burnet and Rufener, 1979), de-
voted to the Geneva 7-eolour photomet-
ric system. The supernova has been
measured 463 t;mes fram Feb. 24, 1987,
to Jan. 17, 1989 (see Burki et al., 1989).
The variations of the Geneva V mag-
nitude with time is shown in Figure 1.
From HJD 2446980 to 2447285, the de-
cline of the V curve was strictly linear,
with a slope in good agreement with the
thermalization of the y radiation result-
ing from the 56CO radioaetivity.
From HJD 2447285 onwards, the ex-
panding envelope began to become
transparent to y radiation whieh eould
thus escape without being thermalized.
The decrease of the V luminosity be-
came steeper than the former linear var-
iation. From this point, however, the
steadily inereasing relative importanee
of the two elose companion stars of SK
-69°202, which are unavoidable with
the present measurement technique,
cannot be negleeted. The later portions
of the eorrected V light curve is shown in
Figure 2 on a larger scale. The eorrec-
tions applied are respeetively +0.001
mag for the first, and +0.065 mag for
the last value in Figure 2. A good
mathematieal deseription of that portion
of the correeted V light curve is obtained
by fitting a parabola over the interval
HJD 2447270 to 2447450 (see Fig. 2).
After HJD 2447450, the luminosity in
the V band started to deerease more
SloWly than the parabolic deseription.
This inflexion of the V light curve, ap-
proximately 600 days after core eol-
lapse, could be interpreted as the effeet
of an additional energy souree, sueh as
the expected central pulsar in interac-
tion with the surrounding material
ejected by the supernova progenitor
(see Arnett, 1988, for the theoretical
light curves of SN 1987 A w;th var;ous
Pulsar luminosities). One must, how-
ever, bear in mind the possible contribu-
tions of other slower deeaying radio-
nuclides such as 57CO, 44Ti and 22Na
Produced during the initial explosive
nucleosynthesis (Nomoto et al., 1989;
Woosley et al., 1989), and which would
affect the light curves in a similar
manner. According to the model calcu-
lated by these authors, the energy sup-
plied by these radioisotopes exceeds
that of 56Co after about 1,200 days
fOliowing core collapse.
Is there a Pulsar in Supernova 1987 A?
Arecent announcement of the dis-
covery of a pulsar in Supernova 1987 A
in the Large Magellanic Gloud has ex-
cited the world-wide astronomical com-
munity. New observations at the La Silla
Observatory by a group of European
from the Max Planck Insti-
tute for Extraterrestrial Physics and the
European Southern Observatory, how-
ever, do not confirm the reality of this
object. More observations are now
needed to settle this important question.
Searching tor the Pulsar
in Supernova 1987 A
Since the explosion of the now fa-
mous supernova in the Large Magellanic
Gloud on February 23, 1987, astronom-
ers have been eagerly waiting for the
emergence of a newborn pulsar. Gurrent
theories predict that the explosion of a
heavy star as a supernova will result in
most of its mass being blown out into
surrounding space, but also that some
of it will be compressed into an ex-
tremely dense and rapidly rotating neu-
tron star at the centre. Such an object
would later manifest its presence in the
supernova by the emission of regular
light pulses (hence the name "pulsar").
Neutron stars measure no more than
10-15 kilometres across, but they
weigh as much as our Sun which is
about 100,000 times as large.
Of half a dozen pulsars known in
supernova remnants, the most famous
are those in the Grab Nebula and the
Vela Nebula. The detection of a pulsar
inside SN 1987 A, the first naked-eye
supernova in nearly four centuries,
1 The group consists 01 Hakki Ögelman, Günther
Hasinger and Wollgang Pietsch (Max-Planck-In-
stitut für Extraterrestrische Physik, Garehing bei
München, F. R. Germany), Christian Gouiffes, Jorge
Melnick, Thomas Augusteijn, Flavio Gutierrez, Pre-
ben Grosb01 and Christian Santini (ESO) and Holger
Pedersen (Iormerly ESO, now Nordic Optical Tele-
scope Scientilic Association).
would provide the definitive confirma-
tion of the creation of pulsars in super-
nova explosions. Extensive searches for
such a pulsar have therefore been made
at some southern observatories since
the explosion was first recorded, almost
exactly two years ago. This is done by
observing the supernova light with a
"rapid" photometer, capable of measur-
ing the light intensity many times each
second. A pulsar would reveal itself by
the presence of brief "flashes" of extra
light, regularly spaced in time.
Immediately after the explosion, the
dense cloud around the supernova did
not allow a look at its centre, but as the
clouds become thinner, light from the
new pulsar should eventually become
visible. Many astronomers have been
waiting for this exciting moment.
First Detection?
On 8 February 1989, a group of
American astronomers
announced the
discovery of a very fast pulsar in SN
1987 A, f1ashing no less than 1969 times
per second. This is referred to as 1969
cycles/second and supposedly corre-
sponds to the number of rotations per
second by the pulsar. No other pulsar
has ever been found to rotate this fast.
These observations were made on
January 18, 1989 at the Gerro Tololo
Interamerican Observatory, situated
about 100 km south of La Silla. Surpris-
ingly, the American group did not see
any pulsations when the observations
were continued 12 days later with
another telescope.
At the European Southern Observato-
ry, the light of the supernova has been
monitored with a special, rapid photo-
meter at the 3.6-m telescope at regular
2 This group is headed by John Middleditch, Los
Alamos National Laboratory; the announcement
was made on Circular 4735 01 the International
Astronomical Union (lAU).
intervals during the past year. The inten-
sity of the supernova light was mea-
sured 1000 times per second, a rate
which was determined from theoretical
considerations about how fast the pre-
dicted pulsar in SN 1987 A might rotate.
However, it is too slow to show varia-
tions at the rate observed at Tololo.
New Observations Fail to Provide
In order to confirm the presence of a
pulsar with the higher pulsation rate, the
ESO instrument was modified immedi-
ately after the announcement by the
American group, so that it can now me-
asure the supernova light up to 10,000
times per second. On February 14 and
15, observations were performed at the
3.6-m telescope during a total of 8
hours. The data tapes were rushed to
the ESO Headquarters in Garching near
Munich. The detailed results of a careful
analysis at the Max Planck Institute for
Extraterrestrial Physics have now been
published on lAU Gircular 4743 (24 Feb-
ruary 1989).
The European team finds that no pul-
sating signal is present in the ESO data
near 1969 cycles/second, to a limit of
1/400Oth of the intensity of the super-
nova light. Nor is there any obvious sig-
nal at any other frequency in the interval
from 1 to 5000 cycles/sec. These obser-
vations therefore do not provide confir-
mation of the presence of a pulsar.
If there is a pulsar in SN 1987 A, then
the absence of observed pulsations in
the measurements obtained after Janu-
ary 18, both by the American and the
European groups, possibly indicate that
the pulsar is being intermittently ob-
scured by dust clouds around the
Further observations will therefore be
needed to definitively demonstrate the
reality of a pulsar in Supernova 1987 A.
(From ESO Press Release 02/89,
24 February 1989.)
Squeezing the Most from the CES
H. BUTCHERand T. SCHOENMAKER, Kapteyn Observatory, Roden, the Netherlands
In the Messenger No. 51 (p. 12-15),
an attempt to use the GES for radioac-
tive chronometry of the Galaxy was de-
scribed. Two instrumental aspects of
the observations were noted as being
not yet under adequate control: (i) The
determination in detail of the instrumen-
tal profile, the wings of which cannot be
determined weil, except at laser line
wavelengths, and the shape of which is
found to depend on the often variable
focus across the array detector format,
and at least sometimes on the signal
level in the detector used. (ii) The mar-
kedly asymmetrical distribution of the
noise from the Reticon detector (caused
by energetic particle detections) for low
flux levels and integrations of an hour
and longer. Treatment of the data in an
optimum way is problematical due to
Figure 1: Flat field correction frame displayed in four segments, starting at lower left and
finishing at upper right. The frame has a mean correction of 1.000, with fringes and sma/! scale
sensitivity variations visible as deviations from 1.0 wherever there is sufficient signal for
meaningful correction.
these effects, and although the CES re-
mains one of the world's best instru-
ments of its type, study of such deficien-
cies can be expected to improve its
Performance significantly.
A satisfactory way of deriving the in-
strumental profile is, to our knowledge,
still lacking, but with the possibility of
lJsing the new double density RCA CCO
On the long camera, it seemed possible
to substantially improve the noise
characteristics of the data. That is, one
can image the spectrum so that each
spectral channel covers several pixels,
and discriminate and eliminate (mostly
Single pixel) particle detections during
data reduction. When Oaniel Hofstadt
On La Silla was approached in late 1987
about when this detector might be avail-
able to long-camera users, he gracious-
Iy promised to help provide it for a CAT/
CES run on 15-22 April 1988 - but only
if a detailed report of the experience
Would be prepared! This article is
adapted from that report.
The CCO we used was ESO CCO
No. 9, an RCA SIO 503 high resolution,
thinned and backside illuminated de-
vice, with 1,024 x 640, 15 micron square
Pixels. We used a data format of 1,030
rows of 61 pixels each, with the spec-
trum recorded along the long dimen-
sion. The noise is quoted at 40 electrons
rms when used in the fast readout
mOde. We decided to use the slow read-
Out mode, however, for which we could
discover no precise noise figure; our
data are consistent with a system noise
of about 35 electrons rms. The dark
current was said to be 2.5 electrons per
pixel per hour, a value which agrees with
our data. The charge transfer efficiency
was reputed to be very good for this
chip. Our spectral lines were all in the
blue region, near 4000 Ä, and we were
pleased to discover that the overall de-
teetion efficiency at this wavelength was
very reasonable. A bad column was
present near the centre of the chip, but
our sub-format could be set up to avoid
it. The spectrum signal showed a rough-
Iy gaussian shape in horizontal profile,
some 5 pixels FWHM for the stellar data
and 8 pixels for the calibrations.
All things considered, this chip seems
to be representative of the higher quality
devices available to astronomers today,
and we were particularly interested to
determine what the limits of its photo-
metrie performance would be.
To try to extract stellar spectra in an
optimum way from the 2-D CCO data,
we have used the following processing
(1) Background subtraction and flat
We use a constant for the
background, because all our dark
frames were beautifully uniform and flat.
We had three kinds of flat field data:
continuum light from the internal cali-
bration lamp, from a rapidly rotating hot
star observed under similar circumstan-
ces to our programme stars, and from
an illuminated patch on the inside of the
dome observed through the telescope.
Unfortunately, the intensity of the dome
flat field was never sufficient in the blue
to allow adequate S/N to be achieved.
The other two sources yielded very simi-
lar results: the spectra, even at 4000 Ä,
show fringing at the 5% level, but luckily
this fringing is rather stable and we
could not really distinguish between in-
ternal lamp data and the stellar spec-
trum for this purpose.
We prepared the flat field correction
by heavily smoothing the continuum
spectrum in the direction of dispersion,
then dividing the raw spectrum by the
smoothed one, to yield a flat correction
picture with data values near 1.0. For
those parts of the frames where the
signal is so low that this procedure
would objectionably amplify the noise,
or where division is by nearly or exactly
zero, we simply replace the ratio value
by exactly 1.000. This correction picture
is then divided into the data frames,
removing high frequency sensitivity vari-
ations and fringes, while leaving the low
frequency echelle order intensity en-
velope unchanged. An example of the
resulting flat field correction frames is
shown in Figure 1. Again the fringes
have peak-to-peak amplitudes of
roughly 10%, and in pi aces are of such
a spatial frequency and orientation that
they could cause significant line
strength errors without obviously being
seen to do so.
(2) Extraction of one dimensional
Because we want to maximize the
final S/N in our extracted spectra while
minimizing the effects of particle detec-
tions and single pixel defects, we
choose to derive a template function in
the direction perpendicular to the spec-
tral dispersion, and fit this function by
iterative least squares at each position
along the spectrum.
The template is derived by averaging
on the order of a hundred rows of data
(each row consisting of 61 pixels cutting
across the spectrum), selected from the
centre region of the format so as to
avoid obviously bad pixels. A typical
template is shown in Figure 2a. This av-
eraging process ensures not only that
the S/N of the template signal is much
higher than that of the data to be fitted,
but also that it represents an average
signal strength in the data at hand, and
hence should suffer minimally from any
photometrie non-linearities present.
Our fit of the non-analytical templates
to the data assumes that three parame-
ters are involved: a (flat) background
level, a multiplicative gain factor for in-
problem columns
tion along the spectra. As have other
workers, we find that the scatter around
the fit of the thorium emission fines is
rather larger than estimated from the
noise in the data. The resulting
wavefength scale is adequate for our
purposes, but it is clear that future cali-
bration facilities for such high quality
spectrometers should incorporate a sol-
id Fabry-Perot etalon, to yield many
sharp lines with precisely known spac-
ings for the dispersion calibration.
We also note that our thorium cali-
bration data show wavelength zero
point jumps during one night of up to
1.9 km/so We do not know the origin of
these jumps, but evidently for radial ve-
locity programmes it is important to take
wavelength calibration spectra at regu-
lar intervals with the CES.
Assigning a temperature to the inter-
nal lamp spectrum of 2,800 K and ratio-
ing it to a black body spectrum at that
temperature yields a correction to the
large scale sensitivity variations from
one end of our spectra to the other
which we estimate are accurate to bet-
ter than 10%. The main uncertainty
arises because of the differing optical
paths of stellar and lamp light in the
instrument, which in principle can give
rise to illumination differences of the op-
tics. At least to first order, however, this
strategy provides a useful correction.
Posl t.ion
Figure 2: (a) Typical signal template, used for fitting to each spectral channel. (b) Trace of data
perpendicular to spectrum in the bottom of an absorption line, showing anomalous columns.
tensity, and a shift in the x coordinate
(i. e., perpendicular to the spectrum). At
each channel (y position) along the
spectrum we fit the template, calculat-
ing the derivative of the template with
respect to the x coordinate using fifth
order polynomial interpolation. The rms
difference between the template and
data is minimized by iterating until the
calculated parameter changes are no
greater than those expected from the a
priori known noise, or until a maximum
number (usually 7) of iterations has been
made. For reasons sketched below, it
has been found necessary to apply a
mask to eliminate some of the data from
the fitting process.
In Figure 3 is shown a processed
spectrum of a 3.6-mag star, trom an
integration of 15 min, and the difference
between it and a second integration of
the same star immediately following.
The signal levels in the two cases differ
by more than 30 %, due to variable see-
ing. It is evident that the fitting proce-
dure automatically normalizes all spec-
Gt.B03 41. 1:l1:l0
corrected) f1at field continuum data to
provide a first correction for the large-
scale sensitivity and echelle order varia-
21. eBG 1. GI:lB
Figure 3: Extracted spectrum from a single, well-exposed spectrum of a solar type dwarf star.
The spectral resolution is 100,000. The strong absorption line at pixel position 720 is Fel
4132 A. The noise trace near zero intensity is the subtraction of two such spectra, which differ
in recorded signal level by 30-40 %.
(iii) Wavelength scale and continuum
Finally, we have used the internal
thorium lamp at the CES to obtain a
wavelength calibration, and have ex-
tracted a spectrum from the (fringe
~ e e . ooe 6ee. eee 7ee. aee 800. eaa
to a common intensity scale (no scal-
Ing in intensity was made here), and that
t,here is no significant detector non-
i1nearity present (which would be mani-
fest as residual differences at the posi-
tions of strong Iines). Furthermore, the
difference of the spectra has an rms
noise very near to the calculated photon
shot noise limit, given the quoted 6.8
electron/ADU calibration.
From these and other results from the
rUn it is clear that apparent S/N ratios of
300 : 1 are achievable with this detector.
ratios may be possible with mul-
tiPle integrations.
The only serious problem found in our
data is shown in Figure 2b. The two
Columns marked there consistently
Show residuals (data-minus-fit) depen-
dent on signal level. We understand
(Sandro O'Odorico, private comm.) that
these RCA chips are known to exhibit
such behaviour - that is, to have occa-
column pairs in which part of the
Signal in one column seems to end up in
the adjacent column, when the signal is
above some threshold level. We have
experimented with trying to fix this prob-
lem by applying an interpolated re-
transfer of signal after the fact, but could
not Convince ourselves that the results
were always reliable.
Gur solution to the problem is to mask
the offending columns of data away,
exclude them from the template
process. We thereby lose some
5% of our signal, whichwe deem an
acceptable loss to guarantee the quality
of the photometry. The data in Figure 3
were processed with these columns
masked away.
One negative side effect results, how-
ever. With two fewer signal columns, the
fitting algorithm now no longer effective-
Iy ignores the single pixel outliers due to
particle detections. We have had to in-
clude a routine, therefore, to test for
pixels more than 5 times the (a priori
known) noise sigma from the fit, and
throw the worst single one out. Although
not particularly elegant, this strategy has
proven very effective in removing parti-
cle detections.
Focus Effects
Our data-minus-fit residual frames are
exquisitely sensitive to focus variations
along the spectra (although the final in-
tegrated intensities at each point should
not bel. We find that the width of the
spectra does vary, being broadest at the
two ends, but that the effect is so small
that we cannot measure it in the widths
of individual emission lines in the cali-
bration spectra. We suppose that a
small tilt of the CCO with respect to the
focal plane of a half degree is sufficient
to give the magnitude of what we see.
From the residual frames we infer that
over about a quarter of the total length
of the recorded data, the instrumental
profile is clearly constant enough for use
in spectral modelling analysis with a sin-
gle model profile. But of course, the
derivation of the appropriate profile re-
mains problematical.
Based on our, admittedly incomplete,
analysis, we feel we can make the
following conclusions.
(i) The double density RCA CCO on
the CES long camera works very weil,
even at 4000 A.
(ii) Its lower noise per pixel compared
to the Reticon, and its registry of each
spectral channel with multiple pixels,
allows particle detections to be discov-
ered and easily removed.
(iii) The expected quantum noise limit
is achieved on single integrations, allow-
ing S/N ratios of several hundred to be
obtained. We have not done tests to
determine whether S/N ratios of 1,000
and greater are possible, by summing
multiple integrations.
(iv) Photometrically unacceptable col-
umns on the detector have been discov-
ered, which should be masked out dur-
ing analysis if results of the highest qual-
ity are to be attained.
(v) Least squares fitting of templates
for data extraction, and probably also
for wavelength calibrations, seems a
good way to determine the length of
spectrum over which the instrumental
profile is sensibly constant in shape.
'Nhat is the Mass-to-Light Ratio of the Old Magellanic
Globular Cluster NGC 1835?
p. DUBATH, M. MAYOR, Observatoire de Geneve, Switzerland, and
p. MAGAIN, Institut d'Astrophysique de Liege, Belgium
1. Richness of the Southern Sky
h We astronomers are lucky: our Galaxy
as two companion galaxies, the Large
and Small Magellanic Clouds, situated
weil above the galactic plane, which
Contain a huge potential of astrophysi-
?al information. For example, concern-
star clusters, the realm of the globu-
ar clusters is much richer and more
Varied in the Magellanic Clouds than in
the Galaxy: rich clusters of all ages are
observed, from the youngest, having
ages of a few tens 10
yr to the oldest
h ' "
1aVlng ages of the order or larger than
yr. In this paper, only old
Magellanic and galactic globular clus-
ters are considered.
From the determinations found in the
literature of the individual masses of the
richest old clusters, a systematic differ-
ence seems to exist between the globu-
lars in our Galaxy and in the Magellanic
Clouds, Magellanic clusters appearing
fighter than galactic clusters. This differ-
ence in mass between old rich
Magellanic and galactic clusters obvi-
ously has direct consequences on the
mass-to-light ratio determination of the
considered clusters, reflecting perhaps
systematic differences in mass function.
This was challenged and discussed re-
cently (Meylan 1988b). A way to resolve
this controversy consists in obtaining
good observational values of the central
vefocity dispersion, by detecting the
very small line broadening present in the
integrated light spectra.
2. Magellanic and Galactic
2. 1 Mage/lanic globular clusters
The method most often used for ob-
taining the total mass of Magellanic
clusters is related to the systemic rota-
tion of the Magellanic Clouds. The ob-
served value of the tidal radius r
of the
cluster is transformed into mass, in a
way similar to the case of galactic open
clusters. It is assumed that the clusters
TABLE 1: The different values of the total mass of NGC 1835 published during these last
15 years.
Year x M
M/Lv r
MG) [G units) [rcl
" .
0.045 0.2 ., . Freeman 1974
" .
0.044 0.12 . .. Chun 1978
" .
0.062 0.17 . .. Chun 1978
" .
0.073 0.18 . .. Elson & Freeman 1985
1985 ... 0.16 0.42 iso Elson & Freeman 1985
1988 1.75 0.39 1.30 iso Meylan 1988b
1988 1.50 0.28 0.94 30 Meylan 1988b
1989 1.25 1.03 3.58 iso present study
1989 1.00 0.81 2.83 30 present study
are in rotation along circular orbits
around the centre of mass of the
Clouds, the old clusters seeming to form
a disk-like system in the LMC. Tidal
masses, particularly for the outer clus-
ters, may be underestimated, if these
clusters are in radial rather than circular
Observational determinations of the
tidal radii by star counts in the outer
parts of Magellanic clusters is pioneer
work of a difficult nature, since the
pollution by Magellanic fjeld stars still
induces uncertainties. Determination of
tidal radius is a difficult task even for the
galactic globular clusters, and the "ob-
served limiting radius" determinations
are rather weak for nearly all of them.
Using the above method, Elson and
Freeman (1985) found for the following
three old LMC clusters, NGC 1835,
NGC 2210, and NGC 2257, total masses
equal to, respectively, M
= 7.3, 1.9,
and 3.7 10
MG, with corresponding
mean M/Lv ratios equal to 0.18, 0.11 and
0.56 (fable 1).
Reasonably good dynamical con-
straints - surface brightness profile and
central value of the velocity dispersion -
have been published so far only for one
Magellanic cluster: NGC 1835. It is only
in the case of this cluster that the deter-
mination of the velocity dispersion (o(Vr)
= 5 km S-1, obtained by Eison and Free-
man (1985) from a Fourier method ap-
plied to integrated light spectra) can be
converted into mass, by using the core
radius and the dimensionless mass de-
rived from the fit of the observed surface
brightness profile to single-mass iso-
tropic King models. Elson and Freeman
(1985) find for NGC 1835, M
= 1.6
MG, with MlL
= 0.42 (fable 1).
All the former mass and mass-to-Iight
ratio determinations concerning NGC
1835 (Freeman 1974, Chun 1978, Elson
and Freeman 1985, and Meylan 1988b)
are displayed in Table 1 and will be dis-
cussed in Section 4 with the results of
the present study.
2.2 Galactic globular clusters
In our Galaxy, only six globular clus-
ters have been studied so far with King-
Michie multi-mass anisotropic dynami-
cal models, consisting of about ten dif-
ferent subpopulations. The observation-
al constraints for such models are the
surface brightness in the central parts
and star counts in the outer regions,
with in addition the high quality radial
velocities for numerous individual
member stars. These six best studied
galactic globular clusters are M3 (Gunn
and 'Griffin 1979), M92 (Lupton et al.
1985), M2 (Pryor et al. 1986), M13 (Lup-
ton et al. 1987), (J) Cen (Meylan 1987),
and 47 Tuc (Meylan 1988a, 1989). Apart
from (J) Cen, the (unique) giant cluster of
the Galaxy (M
= 3.9 10
MG), the total
masses range from 0.4 to 1.1 10
whereas all the mass-to-light ratios are
located between about 2 and 3. The
above values can be considered as typi-
cal of the masses and mass-to-light
ratios of the rich globulars of our Galaxy.
2.3 A difference in M/L
bya factor of
The typical mass of the globular clus-
ters in the Clouds (Iess than 10
MG) is
smaller than the typical mass of the
globular clusters in the Galaxy (greater
than 10
MG)' This difference in mass
between rich galactic and Magellanic
clusters obviously has direct conse-
quences on the mass-to-light ratio: M/Lv
= 0.1-0.5 for those in the Magellanic
Clouds and M/L
= 2.0-3.0 for those in
the Galaxy. ls there a genuine systemic
difference in M/Lv, by nearly a factor of
It is worth emphasizing that the above
question does not only concern globular
clusters. For example, despite the range
of a factor of 1,000 in galaxy luminosi-
ties in the Local Group galaxies, the
globular cluster luminosity distributions
of these galaxies are consistent with
being of the same form in all of them. It
is generally accepted that the globular
cluster population in galaxies was
formed with the same distribution of
globular cluster masses and
luminosities in all galaxies. A clear sys-
tematic difference in M/Lv between
Magellanic and galactic globular cluster
populations would cast doubt as to their
use as secondary distance indicators in
the cosmological distance ladder.
At the present time, there is no defini-
tive answer. We emphasize that only
the rich old globular clusters in the
Magellanic Clouds and in the Galaxy are
considered here. lt is essential to realize
that any comparison between the M/Lv
values of these two populations is so far
strongly hampered by the fact that these
values proceed from different determi-
nation processes. Comparison between
galactic and Magellanic M/L
should be done only between results
coming from the same kind of models
constrained by the same kind of obser-
vational data. The more elaborated
King-Michie dynamical models have
been applied so far only to galactic
globular clusters (with the exception of
NGC 1835), due to the lack of observa-
tional data concerning the Magellanic
This situation is on the verge of
change: if it is still difficult to obtain
velocity dispersion profiles of
Magellanic clusters, it appears noW
feasible to obtain at least the central
value of the velocity dispersion, from
integrated light spectra. The obtaining
of such an essential observational con-
straint, in the case of the old Magellanic
globular NGC 1835 is presented below,
with a discussion of the application of
this new result to a King-Michie model
and the consequences on the M/L
concerning this cluster.
3. Core Velocity Dispersion from
Cross-Correlation Technique
3. 1 Optical and numerical cross-corre-
For more than twenty years, the
cross-correlation spectroscopy has
proven its exceptional efficiency in ra-
dial velocity determination (Griffin 1967,
Baranne, Mayor, and Poncet 1979). The
cross-correlation between a stellar
spectrum and a template allows to con-
dense the radial velocity information,
contained in the stellar spectrum, into
the equivalent of a single spectral "line":
the cross-correlation function. For ex-
ample, with the two CORAVEL spec-
trometers, using simultaneously about
1,500 spectral Iines, a few minutes of
integration provide adetermination of
the radial velocity of a 14th V-magnitude
star with aprecision of about 1 km S-1.
In such spectrometers, the cross-cor-
relation is done optically, but it is easy to
During the integration, a scanning of the
nucleus was done with the entry slit, in
order to cover a zone of 6 x 6 arcsec, so
that the velocity dispersion is based on
the light coming from more than 100
giant stars. Data reduction of the spec-
tra was done with MIDAS. It is worth
mentioning that the cross-correlation
technique allows us to determine accu-
rately a broadening of only a few per
cent of the FWHM of the cross-correla-
tion function (at zero broadening factor,
nique, we obtained two hours of a test
night at the ESO 3.6-m telescope. Spec-
tra were obtained for two old globular
cluster nuclei of the Large Magellanic
Cloud as for two K5 111 comparison
stars. The integration times were two
20-minute exposures for NGC 1978 and
two 30-minute exposures for NGC
1835. Due to the large difference in cen-
tral surface brightness between the two
clusters, the signal-to-noise ratio of the
two cluster spectra are quite different
(Fig. 1). The case of NGC 1978 is espe-
cially interesting for illustrating the
potentialities of the method: the spec-
trum of this cluster has only a signal-to-
noise ratio of about 2.
To increase the sensitivity to line
broadening, we prefer to exclude strong
saturated lines in the correlation pro-
cess and use only small unsaturated
lines selected from the CORAVEL tem-
plate (here using its numerical version).
Only the spectral domain of the
CORAVEL mask between 4400 and
5200 A has been used when cross-
correlating CASPEC spectra with the
numerical mask. In the left half of Fig-
ure 1, 100 Aranges from the spectra of
one comparision star (HD 31871, K5111,
= 9) and from the two clusters NGC
1835 and NGC 1978 are displayed. In
the right half of the same figure the
cross-correlation functions obtained for
these objects are plotted. Due to their
low metallicities, the cross-correlation
functions of the two clusters are much
less contrasted than the cross-correla-
tion functions of the two comparison
stars (the width of the cross-correlation
function does not depend on the
metallicity!). After normalization of the
cross-correlation function of the cluster
NGC 1835 in order to have the same
depth as the cross-correlation function
of the comparison star (Fig. 2), we im-
mediately notice the important broaden-
ing of the cluster cross-correlation func-
tion. 80th comparison stars have been
checked by direct CORAVEL measure-
ments to have an almost zero rotation.
Consequently, the velocity dispersion in
the core of NGC 1835 is immediately
Gv, (NGC 1835) = 10.1 km S·1
concerning line broadening measure-
ments, giving access to rotation through
precise V sin i. Even more, the cross-
correlation function of the integrated
light spectra of globular cluster cores
should allow adetermination of the ve-
locity dispersion of the stars in these
cores. Aresolution of about 20,000 is
needed to have access through cross-
correlation spectroscopy to Doppler
broadenings of a few km S-1, but a low
signal-to-noise ratio (as low as 2-3) will
be admissible. A similar approach was
already used by G. IIlingworth in 1976,
to determine the velocity dispersion in
the nucleus of a few galactic globular
clusters, using Fourier transforms of in-
tegrated light spectra (on photographic
3.2 Results trom observations during a
In order to investigate the possibilities
of CASPEC spectra applied to this tech-
4920 4940 4960 4980 5000 240 260 280 300
Wavelenglh (A) Velocily (Km/s)
1: Left: for the comparison star HO 31871 (K5111, m
= 9), and for the two o/d LMC
clusters NGC 1835 and NGC 1978, 100A ranges from the spectra obtained with
S;SPEC (Cassegrain ESO Echelle Spectrograph mounted on the ESO 3.6-m te/escope at La
) are disp/ayed. Right: numerica/ cross-corre/ation functions for the same three objects;
Ue to their /ow metallicities, the cross-corre/ation functions of the two clusters are much /ess
COntrasted than the cross-corre/ation function of the comparison star.
Visualize obtaining the same result from
a numerical correlation between regis-
tered spectra. Such radial velocity mea-
Surements are obtained by D. Latham
and his collaborators, from numer-
Ical cross-correlation of registered
spectra having a wavelength range of
?O A. A numerical cross-correlation, us-
Ing CASPEC spectra, would take advan-
tage of a much larger spectral range
(Iarger than 1000 A). The numerical
cross-correlation gives apriori a notice-
gain: the scanning required to build
optlcally the cross-correlation function
on the telescope is no longer necessary,
providing an immediate gain of about
2.5 magnitudes. In addition, we can ex-
pect some further gain due to the high
quantum efficiency of CCDs as com-
Pared to photomultipliers. Unfortunately
t,he readout noise of the CCD is still the
Iimiting factor.
. If the cross-correlation spectroscopy
IS .well adapted to radial velocity deter-
mlnations, it shows the same efficiency
4. NGC 1835 M/Lvtrom King-
Michie Model
Il 11(.]
Figure 3: NGC 1835: /ogarithm of the ob-
served (dots) and mode/ surface brightnesS
profi/es as a function of the /ogarithm of the
..8 1
mass of NGC 1835 are larger than the
best previous determinations (not using
King-Michie models) by about a factor
of ten, giving to this old LMC globular
cluster aM/Lv ratio similar to those ob-
tained for galactic globular clusters
(Table 1). Consequently, the method
based on the assumption of circular or-
bit (due to rotation) and on the transfor-
mation of the tidal radius into mass
should be used only with great care in
the case of Magellanic clusters. The
systematic difference observed be-
tween the typical mass of rich globular
clusters in the Galaxy (greater than
MG) and in the Magellanic Clouds
(Iess than 10
MG) could be a simple
artifact, a direct consequence of the
idiosyncrasies of the different methods
In conclusion, when the same kind of
dynamical models (King-Michie) con-
strained by the same kind of observa-
tions (surface brightness profile and
central value of the velocity dispersion)
are applied to an old rich MagelianiC
globular cluster, viz. NGC 1835, the re-
sults seem similar to those obtained in
the case of 47 Tucanae. Consequently,
the rich old globular clusters in the
Magellanic Clouds could be quite simila
(in mass and M/L
) to the rich globul
clusters in the Galaxy. The presen
study will be extended to as many
Magellanic clusters as possible.
Baranne, A., Mayor, M., and Poncet, J. L.
1979, Vistas in Astronomy, 23, 279.
Ghun, M.S. 1978, Astron. J., 83, 1062.
Elson, R.A.W., and Freeman, K.G. 1985, As-
trophys. J., 288, 521.
Freeman, K.G. 1974 ESO/SRG/GERN Gon-
ference on Research Prog. for the NeW
Large Telesc. Geneva p. 177.
Griffin, R. F. 1967, Astrophys. J., 148, 465.
Gunn, J. E., and Griffin, R. F. 1979, Astron. J.,
Lupton, R., Gunn, J. E., and Griffin, R. F. 1985
. Oynamics of Stars Clusters lAU Symp·
Within the large range of values of x
investigated (from 0.0 to 3.5 by steps of
0.25), only the values between 1.00 and
1.75 provide models able to fit the ob-
servations. Only models with small frac-
tions of stellar remnants (neutron stars
and white dwarfs) fit the observations:
the fraction of the total mass in the form
of neutron stars varies from 0.0 to 4%,
whereas the fraction of the total mass in
the form of white dwarfs varies from 9 to
26%, depending on the model.
From a structuraf point of view, NGC
1835 appears rather concentrated, with
values of the concentration parameter c
= log (rt/rc) ranging from 1.81 to 2.24
(similar to 47 Tuc). It is worth mentioning
that the size of NGC 1835 (rt = 50 pc) is
quite comparable with the size of w Cen
and 47 Tuc.
Oepending on the model, the values
obtained for the total mass of the cluster
range from 0.70 to 1.55 10
MG, with a
representative mean total mass (M
) =
1.0 10
MG' From Table 1, we see that
the best results obtained by transform-
ing the tidal radius rt into mass (under
the assumption of systemic rotation of
the old globular cluster system) are
smaller than the results from King-
Michie models by about a factor of ten.
The half-mass relaxation time and the
central relaxation time are of the order of
10 Gyr, and 10
yr respectively, allowing
a large fraction of the central parts of the
cluster to have been relaxed. Both iso-
tropic and anisotropic models are suc-
cessfully fitted to NGC 1835. Areal ve-
locity dispersion profile, instead of pres-
ently only the central value, would allow
perhaps a better evaluation of the quan-
tity of anisotropy.
The central surface brightness 110 var-
ies from 6.53 to 6.72 mag/arcmin
with a
representative mean value (Ilo) =
6.66 mag/arcmin
• The integrated visual
magnitude V
varies from 9.85 to 9.91
mag with a representative mean value
) = 9.88 mag. This last model value is
situated between the observed values:
= 9.48 (Chun 1978), V
= 9.52 (Elson
and Freeman 1985), and V
= 10.13 (van
den Bergh 1981). The global mass-to-
light ratio M/L
varies from 2.50 to
5.21,with a representative value (M/L
= 3.4, whereas the centraI mass-to-light
ratio (M/Lv)o varies from 1.93 to 2.63.
These values of the mass-to-light ratio
are the direct consequences of the cen-
tral value of the velocity dispersion, and
depend also mainly on the mass func-
ti on exponent x and on the quantity of
anisotropy of the velocity dispersion.
5. A Universal M/L
tor Old
Globular Clusters?
The present results (using King-
Michie models) concerning the total
40 -20 0 20
Velocily (Km/s)
4. 1 Observational constraints
With this new determination for the
velocity dispersion in the core of the old
LMC globular cluster, a0J,) = 10.1 km
S-1, it is possible to constrain King-
Michie dynamical models. The other ob-
servational constraint, viz. the surface
brightness profile, is a composite profile
- namely CCO surface brightness pho-
tometry in the very central parts, centred
aperture photometry and drift scan
measures in the central and intermedi-
ate parts, and star counts in the outer
parts - obtained by mixing Eison and
Freeman (1985) and Mateo (1987) data
(al ready used in Meylan 1988b).
Figure 3 displays the surface bright-
ness profiles, as a function of the radius,
of the observations and of one of the 8
best models (Iowest reduced ob-
tained from a grid of about 400 models.
The model profile, integrated afong the
line-of-sight, concerns only the sub-
population containing giants, subgiants,
and stars at the top of the main se-
quence, Le. all the stars emitting most
of the light of the cluster. The residuals
between observations (dots) and model
(continuous line) are also displayed in
the lower part of the same figure.
the FWHM of the CASPEC cross-corre-
lation function is equivalent to about
18 km S-1, the pixel size being 9 km S-1).
Figure 2: Normafized cross-corre/ation func-
tions of the cluster NGC 1835 (triang/es) and
of the comparison star HO 31871 (dots); the
continuous fines are the corresponding fitted
gaussians; the important broadening of the
cluster cross-corre/ation function is con-
spicuous and allows an immediate determi-
nation of the ve/ocity dispersion in the core of
NGC 1835: a(V,) = 10.1 km S-I.
! .......
4.2 Mass function exponent, total mass,
and mass-to-light ratio
The models are calculated by using
the same mass function exponent x
(dN oc m-Xdlog(m)) for the entire range
in stellar mass, i. e. from 0.1 to 100 MG'
113, Princeton 1984, Edit. J. Goodman
and P. Hut, p. 19.
LUpton, R., Gunn, J. E., and Griffin, R. F.
1987, Astran. J, 93, 1114.
Mateo, M. 1987 Ph. D. Thesis University of
Meylan, G. 1987, Astron. Astrophys., 184,
Meylan, G. 1988a, Astrol? Astrophys., 191,
Meylan, G. 1988b, Astrap/lys. J. 331, 718.
Meylan, G. 1989, Astron. Astrophys. in press.
Pryor, C., McClure, R.D., Fleteher, J. M.,
Hartwiek, F.D.A., and Kormendy, J. 1986,
Astron. J, 91, 546.
High-tech Telescope on Top of Mount Wendelstein
H. BARWIG, Universitäts-Sternwarte München, F. R. Germany
Figure 1: Ready for observation: fhe new fe/escope on Mount Wende/stein.
A prominent feature of the silhouette
of the Alps, tl1e Wendelstein Mountain,
can be seen from Garching on clear
days. Extremely transparent skies
sometimes even allow one to recognize
On top of that mountain the domes of
lhe Observatory of the University of
Munich (USM). Only 75 km away, after
One hour's ride, this site can easily be
reached via cable car or by means of a
famous 75-year-old cog railroad. Final
access to the very top is achieved by an
elevator climbing up 114 m within the
This Observatory had already a long
tradition in solar research. First observa-
tions were started in 1940 (however,
near the end of World War 11 they merely
served the military to forecast radio dis-
turbance caused by solar activity). La-
ter, observations of solar prominences
and of the solar corona began. After
Inlegration of the Observatory into the
Vvorld-wide activities during the "Inter-
national Geophysical Year" a 20-cm
Zeiss Coude Coronograph was in-
Slalled; it was used since 1963 to pro-
v1de data for international solar re-
About 20 years later, when corona
Observations were reduced to only a few
~ a y s a year due to increased air poilu-
Ion, these activities had to be stopped:
It Vvas high time to think about the future
of the Wendelstein Observatory. On the
other hand, a statistical evaluation of
Vveather data obtained over 7 years
Clearly indicated favourable conditions,
at least for stellar observations, on
about 145 days a year and fairly good
Seelng quality as weil.
Thus in 1983 a new concept was
elaborated to adapt the facilities to
nlght-time observations and to move
lhe solar instruments to a site on the
Canary Islands. After many struggles for
lhe required financial support, finally the
green light was given to the telescope
bUllder OFM Engineering in Colorado to
Provide an 90-cm fork-mounted Ritch-
eY-Cretien telescope and auxiliary
eqUipment, including a grating spectro-
graph and a CCO camera with an image
data analysis system.
Construction work for the new dome
bUilding (Fig. 1) started in autumn 1987.
The telescope itself was delivered in
November 1988, packed into a huge
container that arrived at the bottom of
Wendelstein after a journey of about
9,000 km. Helicopter flights scheduled
for transportation to the mountain top
had to be cancelled due to heavy snows
and wind velocities up to 200 km/ho
Therefore all of the equipment had to be
brought up to the top using the cog
railroad and two elevators.
The mechanical and optical setup
was achieved within the scheduled time,
although some problems were encoun-
tered in getting the telescope drive sys-
tem to work because of strong signals
from the nearby radio station. Extensive
sllielding measures finally solved this
problem. "First light" through the tele-
scope could be announced on January
18, 1989 when radiation reflected from
Jupiter first passed through the new in-
However, long before this date our
request to ESO concerning support of
optical tests on the planned telescope
had been answered positively - another
example of the fruitful cooperation and
of the mutual exchange of experience
we had enjoyed with ESO in the pas!.
For test purposes the ESO-Shack-Hart-
mann camera "ANTARES" could be
made available just before this instru-
ment had to be shipped to Chile for
application as wavefront sensor at the
ND. The Hartmann exposures will serve
to perform the final optical alignments
and to determine the overall optical
quality. Prior to the completion of these
tasks visual observations have already
yielded a surprising result: The atmo-
spheric conditions on Wendelstein dur-
ing several test nights yielded seeing
disks weil below 1" (Craters on the
Moon about one arcsec of diameter still
showed shadow structures in their cen-
The probability of having such favour-
able atmospheric conditions as weil as
the number of useful nights promised by
weather statistics seem to justify the
installation of such a high-tech tele-
scope. Its performance differs in several
respects from those of conventional in-
struments of comparable size: The thin
primary mirror of accordingly low heat
capacity is supported by an airbag. Its
pressure is adjusted according to the
mirror's weight on three hard points.
This device guarantees preservation of
the mirror's shape and therefore optimal
image quality regardless of the tele-
scope position. Furthermore the light-
weight mirror allows a very stiff mount-
Figure 2: The Multichannel Photometer
mounted at the 0.8-m OFM telescope.
ing of low inertia to be used, which
makes possible a slew speed of 4 de-
grees/sec. Computer control provides
corrections for precession, aberration,
nutation, refraction, azimuth and eleva-
tion misalignment, mechanical and opti-
cal non-perpendicularities and flexure.
These corrections typically allow a
pointing accuracy of better than 10 arc-
sec RMS to be achieved. Back-up in-
strumentation weighing up to 100 kg
can be mounted at the instrument
rotator. (Highly sensitive detectors
which may be influenced by radio noise
can be protected by means of a shield-
ing box for RF radiation.) The telescope
is operated from aseparate control
room. Two terminals respectively allow
one to command the telescope motion
via a menu and to display all relevant
status parameters. An automatic dome
drive and an autoguider that follows
stars brighter than 15 mag provide an
almost automatic observation mode.
The first astronomical instrument
mounted (Fig. 2) was the High Speed
Multichannel Photometer MCCP (cf. The
Messenger No. 48, p. 29) developed at
the USM which allows simultaneous
UBVRI measurements of object, com-
parison and sky, thus compensating for
changing atmospheric transparency.
(This instrument has already been suc-
cessfully operated at the ESO 1-m,
2.2-m and 3.6-m telescopes.) High-
speed light curves of the dwarf nova U
Gem were the first astronomical data
recorded on Wendelstein. Application of
the remotely controlled spectrograph
with a Reticon detector and the CCO
camera for narrow band imaging are
planned in the near future.
In principle the new telescope should
serve the following specific purposes:
- Thorough training of students and ob-
servers as preparation for observing
runs at large telescopes.
- Performance of test runs in the
course of instrument and detector de-
- Execution of observing programmes
which are so time-consuming that ob-
serving time at other sites will not be
given to visiting astronomers - or pro-
grammes which concern singular events
like nova and supernova eruptions or
the appearance of comets. All of these
tasks are favoured by the easy access
to our observatory, as weil as by the
existence of a proper infrastructure.
In addition, astronomical observa-
tions simultaneous with those on
satellite-borne telescopes, and cooper-
ations with international observing cam-
paigns are planned. Finally, a certain
amount of telescope time will be
granted to visiting astronomers to alioW
them to take a look through the blue
Bavarian sky.
We would like to acknowledge the
excellent support we encountered by
people of the ESO Optical Group, in
particular by P. Giordano, F. Merkle, L.
Noethe and R. Wilson.
Spatial Resolution Imaging of the Radio Source 3 C 255
Figure 1: A 90-minute eeo frame (V) ob-
tained with the 2.2-m telescope at La Silla,
showing the distant cluster of galaxies
around the radio source 3e 255. The mag-
nitude of the optical image of the source is
- 23. The foreground cluster in the upper part
of the frame is at redshift z - 0.2. North is up
and East is to the right.
The optical identifications of the 3CR
sources are now nearly comp/ete (Spin-
rad et al. 1988), and only 7 are lacking
The radio source 3 C 255 is one of
these remaining sources, and the deep
image presented here was taken to
search for a possible underlying distant
cluster. Three CCO exposures in V and
R were taken with the 2.2-m ESO-MPI
telescope on La Silla under moderately
good seeing conditions (1.2 to 1.5 arc-
sec). Adding these images gives a total
exposure time of 2 h. The visible coun-
terpart of the radio source is the central
object of the frame. There are three
other objects, W, E, and S, located at
about 7 arcsec from it.
Spinrad et al. noticed that the central
object is elongated, as most distant
3CR sources are. On this CCO frame we
resolve the elongated region of the
source into at least three components
(possibly five), the average distance be-
tween the components being about
2 arcsec.
The source itself may have four unre-
solved components (on 1 arcsec scale)
~ : . ~ .. ~ \ . ~
Ir· '
~ ....
." ...
. . ..
c • • ,T .•
" .-..
or extensions, namely the central objeet,
two extensions in the northern direction
and one towards the south.
Figure 2: A spatially resolved image of t!18
elongated radio source 3e 255 obtained with
the 2.2-m telescope on La Silla. Exposur8
time: 3600 s + 1800 s in V(seeing 1.2 arcse
+ 1800 s in R (seeing 1.5 arcsec). The optica
Iy identified source is resolved into a brig/
object and at least 3 fainter components.
North is up and East is to the right.
The simplest explanation is that we
have found a distant aggregate of galax-
ies, the radio source being the first rank-
ed object. Its magnitude suggests a red-
shift of Z 2:: 0.6.
Clusters of galaxies have been found
LJp to Z = 0.92 (Gunn et al. 1986) and it
has been suggested that some of the
rnost distant 3CR galaxies, near z ~ 1.8
rnight be cluster cores in the process of
formation (e. g. 3C 326.1, Mac Garthy et
al. 1987; 3C 294, Spinrad et al. 1988).
The dynamical time of galaxy clusters
is of the order of the Hubble time. It
follows that clusters should be dynami-
cally young and may have had different
properties at z 2:: 1. It is also by no way
obvious that the universe contains many
clusters at high redshift. Thus it would
be important to have a spectrum of this
Several distant 3CR elongated galax-
ies have been found to have a complex,
probably multiple structure (Le Fevre
and Hammer, 1988). These authors
have proposed that some of the 3CR
distant galaxies might be affected by
gravitational amplification or lensing by
foreground galaxies, an hypothesis that
mayaiso apply to this new case.
Gunn et al., 1986, Ap. J., 306, 30.
Le Fevre and Hammer, 1988, Ap. J. Letters,
Mac Carthy et al., 1987, Ap. J. Letters, 319,
Spinrad et al., 1988, A. J., 96, 836.
The Spectacular Binary System PG 1550+131
PG 1550+ 131 was known as a faint,
very blue object 0! =16.8, U-B =-1.2) in
the constellation Ophiuchus. Its optical
Spectrum showed the Balmer lines and
the Balmer jump in emission. Scarce
Photometric data indicated large am-
Plitude variations. So it seemed to be a
relatively uninteresting member of the
cataclysmic variables, not deserving fur-
ther detailed observational attention.
Nevertheless, it was included in a pro-
gramme to search for eclipsing, faint
Cataclysmic variables and aiming at the
1:20 ur
1:31 ur
determination of primary masses in
such systems. The observations were
performed using the CGO camera at the
Oanish 1.5-m telescope at La Silla. PG
1550+ 131 turned then out to be the
most spectacular eclipsing system
found during this survey.
On 1988 July 2 at about UT 0 : 28,
after monitoring this object for approxi-
mately one hour, it suddenly dis-
appeared almost completely from the
frame and reappeared after some seven
minutes, indicating the occurrence of a
1:24 ur
1:35 ur
short, very deep eclipse. On-line reduc-
tions revealed a regular, sinusoidal vari-
ation of the brightness, quite different
from that of common cataclysmic var-
iables and it soon became clear that the
next occultation would occur about
three hours later. In fact, this following
eclipse and another one during the next
night could be observed at the expected
The photo shows a sequence of five
CGO images covering the third eclipse
monitored on 1988 July 3. At UT 1 : 20
1:28 ur
Eclipse on July 3,1988
the eclipse has not yet begun; at 1 : 24
the bright star in the system has
dimmed notably; at 1 : 28 it is complete-
Iy eclipsed; it weakly reappears at 1 : 31
and reaches its normal brightness at
1 : 35. The exposure time was three
minutes for all frames. The main spectral
response is in the red region since no
filter was used.
Folding the data with the (orbital)
period of 187 minutes yielded a smooth,
sine-shaped (full amplitude - 0.6 mag)
light curve outside eclipse which occurs
half aperiod after maximum light. Its
depth is at least 4.8 mag and the time
between first and last contact amounts
to about 12 minutes. In fact it must be
deeper and narrower since due to the
relatively long integration time its true
shape is not resolved. Thus PG
1550+131 exhibits eclipses which are
among the deepest if not actually the
deepest ever recorded for a binary. Two
EFOSC spectra obtained near maxi-
mum respectively minimum light dem-
onstrate that the Balmer emissions (indi-
catively superimposed on absorptions)
disappear near minimum light, leaving
the Balmer series in absorption.
These results show that PG
1550+131 is a precataclysmic rather
than a cataclysmic binary. It consists of
a small hot degenerate object (very
probably a white dwarf with T =
18,000 K) and a late main-sequence star
(T = 3,000 K) which does not yet com-
pletely fill its roche lobe to enable mass
transfer to the compact object typical
for cataclysmic binaries. The compact
object is heating up the facing side of
the companion to about 6,000 K, thus
producing the sinusoidal shape of the
light curve during orbital revolution. The
heated hemisphere is also responsible
for the emission lines which cannot be
seen near the eclipse of the compaet
Only very few eclipsing systems in
this transitory phase are known so far
which allow an accurate determination
of the basic parameters that in turn
place constraints on those for cataclys-
mic variables. A detailed study of PG
1550+131 may therefore contribute to
our knowledge of both the precataclys-
mic as weil as the cataclysmic states of
binary evolution.
R. HAEFNER, Universitäts-Sternwarte,
München, F. R. Germany
Report on the Last Observing Run of Multiobject
Spectroscopy: OPTOPUS is Alive and Kicking
L. GUZZO, Osservatorio Astronomico di Brera, Mi/ano, Italy
of a number of them, mainly due to
imperfect centring of the microienses at
their input ends. As the interest in using
500 400 300
~ 70
which must be considered as unsatis-
factory: the poor blue-UV transmission
of the fibres and the reduced efficiency
o 60
~ 50
..... 40
600 700
Figure 1: Total transmission (including reflection losses) of a 3-m polymicro fibre used in the
new OPTOPUS head.
1. Fibre Multiobject Spectroscopy
The ESO fibre facility for multiobject
spectroscopy at the Cassegrain focus of
the 3.6-m telescope, OPTOPUS, has
been operating since March 1985. A
comp/ete description of the system is
given in the ESO Operating Manual
No. 6. It is possible to observe with
OPTOPUS a maximum of 52 objects
distributed over a field of 33'. Special
aperture plates are prepared in advance
of the observations in the ESO work-
shop from accurate Cl and ö coordinates
of the selected objects. These plates are
eventually mounted at the telescope
and the fibres are manually inserted in
the apertures; at the other end they form
the entrance slit of a CCO spectrograph.
In the last four years OPTOPUS has
been used in 33 nights for 14 different
programmes. It has always operated
with high reliability, collecting some
6,000 spectra. The limiting magnitudes
at a resolving power of about 500 in the
visual are about 18.5 for galaxies and 20
for quasars in a two-hour exposure.
While these limits allow useful work for a
large number of programmes, there are
two aspects of the present system
for extensive surveys of galaxy redshifts
at faint magnitudes (see Guzzo and
Tarenghi, ESO Internal Report, October
1987). The use of an efficient fibre spec-
trograph would reduce by a factor of
20-30 the telescope time normally
necessary for such a kind of work, while
allowing on the other hand to study the
galaxy distribution to much fainter limit-
ing magnitudes than those of the deep-
est existing wide-angle surveys (which
are presently limited to m < 15.5, see
e.g. Geiler et al., 1988 and Giovanelli et
al., 1988). Indeed, to have 30-50 galax-
ies in the OPTOPUS field, one has to
select a sampIe limited to magnitudes
as faint as bj < 18.5-19 (where the
g: 8
'" '"
,.< ,.<
>- >-
"" ""
'" ,.<
5000 5500 6000
'" '"
'" '"
,.< ,.<
>- >-
1=1>- >- >-
"" ""
" ""
'" «
;;; 200
'" «
_ 300
'" «
: 200
Figure 2: Spectra ot two galaxies trom the average ot two 3D-minute OPTOPUS exposures (a:
bj =19.5, z =0.24; b: bj =18.4, Z =0.12). The magnitudes are COSMOS derived values tor the
entire galaxy trom ESO-SRC J. survey plate. The sky has not been subtracted. The redshift is
estimated trom the Hand K fines.
the other at the MT (Parry and Gray,
1986). ESO is currently discussing the
procurement of an analogous facility
through ajoint project with the Observa-
tory of Meudon. This will take time, how-
ever, and ESO will continue to offer the
aperture plate version for the next three
3. The Observations in
September 1988: probing the
Feasibility of a Faint Galaxy
Redshift Survey
Two of us had obtained 4 nights of
observing time with OPTOPUS to inves-
tigate the possibility to use the system
the facility is actually increasing (see
e.g., OAEe Workshop, 1988), ESO
started a programme to improve the
weak points while keeping the basic
2. The New Configuration of
the OPTOPUS Head and the
Future Steps to Improve of the
Extensive tests on the efficiency and
focal ratio degradation of different com-
mercial fibres were carried out in the
ESO laboratory (Avila and O'Odorico
1988, Avila 1988). The fibre selected for
the new OPTOPUS head was an all sili-
ca, wet (high contamination by the OH
radical to enhance the UV transmission)
fibre from Polymicro. The transmission
of 3 metre of this fibre is shown in Fig-
ure 1. A prototype head based on 31
fibres was prepared in the ESO fibre
laboratory and in the workshop in the
first half of 1988. At their input ends, the
fibres were polished and mounted in
connectors which fit in the standard
aparture on the OPTOPUS plates. In this
way the fibres are placed "naked", that
is without microienses, in the focal plane
of the telescope. With a core of 320 Ilm,
they subtend 2.3 are sec on the sky and
since they preserve weil the beam focal
ratio, they can still be used with the F/8
collimator of the ESO B & e spectro-
graph. From laboratory measurements
of the average efficiency of the fibres in
the old head and in the new prototype,
we could predict an improvement in the
efficiency by a factor of two. This was
essentially confirmed by the observa-
tions (see below) which indicate a gain
of about one magnitude with respect to
previous results.
By the end of 1989, ESO plans to
implement a new head with 50 fibres to
be coupled with a new F/6 dioptrie col-
limator now on order. In the final config-
uration, the upgraded OPTOPUS will be
a highly efficient instrument. Further im-
provements can be expected when
eeos with high efficiency and lower
read-out-noise become available.
The weak point or if you want the
bottleneck of the facility will then rest in
the preparation of the aperture plates in
the ESO workshop. It is ademanding
operation in terms of manpower and
ESO is forced to set an upper limit to the
number of plates which can be prepared
for a given observing run. The way to
solve this problem is by building a sys-
tem in which the fibres are positioned at
the location of the targets in the focal
plane of the telescope by a mechanical
device under computer control. Two
systems of this type have recently been
put in operation: one at the Steward
Observatory (Hili and Lesser, 1988) and
Optics and Grisms of EFOSC2
magnitudes are measured from plates
of the ESO-SRC sky survey, J colour): at
these magnitudes the mean density of
objects over the field matches the total
number of fibres and the instrument can
then be used at its maximum efficiency.
It was a clear conclusion of the feasibili-
ty study that it would be possible to
carry out the survey to such a depth at a
rate of about 1.5 square degrees per
night. This means that, for example, to
survey a strip of 100 x 1 square degrees
would take about 70 nights of 3.6-m
telescope time. Here we are not going to
enter into details neither on the best way
to conceive such a survey, nor about the
fundamental importance that its results
would have for the study of the topology
of the Universe and consequently for the
theory of galaxy and large-scale struc-
ture formation. For an excellent discus-
sion of these points see the review arti-
eie by Rood (1989).
Before starting the run, it was planned
to use the new prototype head just for a
couple of nights for testing, and then
shift back to the old version for the rest
of the programme. After the first spectra
were obtained, the improvement in effi-
ciency was so evident that it was de-
cided to go on with it for the whole run.
The ESO grating no. 15 was used to-
gether with the f/8 collimator and f/1.9
blue camera of the B & C spectrograph.
This combination gives a dispersion of
170 Älmm and aresolution (FWHM)
around 10 A. This dispersion is quite
appropriate to study the distribution and
rms velocities of galaxies, for which it is
necessary to keep rms errors on cz to
less than 50 km/sec. We plan to mea-
sure eventually the redshifts using
cross-correlation with template spectra
(generally of a bright galaxy or a K-type
star): this technique permits to reach
rms errors on z between 1/25 and 1/10
of the nominal resolution, depending on
the S/N ratio of the spectrum. This cor-
responds in our case to errors in cz
between 25 and 60 km/sec.
The OPTOPUS starpiates were pre-
pared in Garehing in June 1988. The
galaxy magnitude limited sampies were
extracted from a subset of the Edin-
burgh-Ourham Southern Galaxy Cata-
logue (EOSGC) under completion in
Edinburgh (Heydon-Oumbleton et al.,
1989), kindly provided by the authors.
We observed three areas of 1, 3 and 0.2
square degrees around the South
Galactic Pole, all at the same declination
and with right ascensions around 22 h,
00 hand 03 h, respectively. The largest
region is centred on the rich cluster
Klemola 44. The exposure times were
usually of 60 minutes divided into two
exposures for optimal cosmic-ray elimi-
nation. Three or four fibres were used
during each exposure to monitor the sky
spectrum, while a shorter sky exposure
was observed right after each object
exposure, offsetting the telescope 1
minute north. In Figure 2a, b we present
two partly reduced spectra of faint
galaxies in the field of Klemola 44. The
sky has not yet been subtracted but its
contribution in these spectra is mainly
confined to the emission lines. An esti-
mate of the accuracy with which the sky
can be subtracted will have to wait for a
complete reduction of our data.
The average S/N ratio in the blue-
visual continuum of the spectra is
around 25: with such a value the rms
error on the redshift cross-correlation
measurement is expected to be better
than 30 km/sec. Therefore it appears
that OPTOPUS with the new prototype
head is al ready a weil suited instrument
for redshift surveys down to at least bj =
These magnitudes are integral values
for the entire galaxies which extend be-
yond the finite aperture of the fibres. The
real fluxes collected by OPTOPUS de-
pend on the surface brightness distribu-
tion of the objects. For point-like sour-
ces like asos it will certainly be possi-
ble to reach fainter limits. If we consider
also that a further improvement (of the
order of 10%) will come from the intro-
duction of the F/6 collimator, we can
foresee that the updated version of the
OPTOPUS facility will really be one of
the most efficient facilities of this kind.
It is also interesting to relate about the
"mechanical" efficiency of the system,
i. e. the rate of fields observable per
night. We observed with 25 independent
starplates during 5 actual nights of ob-
servations. This can be considered as
an upper limit, as the level of technical
support to this test run was particularly
Anormal changeover of the starplate
takes around 30 minutes. A complete
exposure run of a starplate, including
mounting, alignment, calibrations, flats,
30 min sky exposure and 1 h scientific
exposure, takes around 130 minutes.
An OPTOPOS run at La Silla is a de-
manding task because of the many op-
erations the observer is involved in: the
In the Messenger No. 52 the con-
struction of a second EFOSC for the
2.2-m telescope was announced.
Mechanically, EFOSC2 is virtually a
copy of EFOSC but comparison of the
optical data (see Table 1) shows consid-
erable differences which for some pro-
grammes will make EFOSC2 the pre-
checking of the starplates, the mounting
of them in the special adaptor, the posi-
tioning of the fibres in the holes, the
alignment on the field in the sky and the
monitoring of the resulting CCO spectra.
All these operations have to be done
accurately but quickly if you want to
maximize your observing time. At times
the work becomes frantic and the ob-
server is usually exhausted at the end of
a run. However, the effort is weil re-
warded, if one considers that we col-
lected some 700 galaxy spectra in 5
nights! A similar number of spectra with
single-object spectroscopy would have
required around 50 nights!
In conclusion, it can be said that the
test run with the new OPTOPUS pro-
totype head has been very successful.
First, it has shown that the partial mod-
ifications introduced in the instrument
so far have already improved its efficien-
cy. On the other hand, it has also clearly
demonstrated that a large-scale survey
of galaxy redshifts down to magnitudes
as faint as 19 is possible, not only in the
dreams of cosmologists, but is real/Y
feasible now, with the instruments we
al ready have.
Avila, G. and D'Odorico, S.: 1988, Proceed"
ings of the ESO Conference on Very Large
Telescopes and their Instrumentation,
M. H. Ulrich editor, p. 1121.
Avila, G.: 1988, ASP Conference Series, S.G.
Barden editor, Vo!. 3, p. 63.
DAEC Workshop, 1988, "Large Scale Struc"
tures: Observations and Instrumentation",
Proceedings edited by C. Balkowski and S.
Geiler, M.J., de Lapparent, V., and Huchra,
J., 1988: in "Large Scale Structures in the
Universe", lAU Symp. 130, J. Audouze and
A. Szalay eds., Dordrecht, Reide!.
Giovanelli, R., Haynes, M. P., 1988: in "Larg
Scale Structures in the Universe", lAU
Symp. 130, J. Audouze and A. Szalay eds.,
Dordrecht, Reide!.
Heydon-Dumbleton, N.H., Collins, G.A., and
MacGillivray, H. T., 1989: M.N.R.A.S., sub"
Hili, J. M., Lasser, M. P.: 1986, SPIE627, 303.
Parry, I.A., Gray, P.M.: 1986, SPIE627, 118.
Rood, H.J., 1988: Ann. Rev. Astron. Astro"
phys. 26, 245.
ferred instrument, even if it is mounted
on a smaller telescope.
The optics were completed by the
end of 1988 and measured in the ESO
optical laboratory. The transmission of
the optics is shown in Figure 1. Com-
pared with EFOSC, the transmission has
been improved by a reduction in the
EFOSC/3.6-m EFOSC2/2.2-m
Camera fjeld (mm) 10 x 15 25 x 25 max.
Fjeld size (arcmin.) 3.6 x 5.8 8x 8max.
Wavelength range (A) 3600-10000 3300 x 10000
Camera focal length (mm)
Piate scale at CCO (/-Im/arcsec)
45 52
Dispersion (Nmm) grism
up to 120
up to 70
echelle 55 28
Resolution with 1" slit up to 2200
up to 3500
Grism #
Ois- Central
Wavelength range Blaze
Blaze abs.
persion wavelength
with TH 1024 x 1024 wavelength efficiency
(Nmm) (A)
1 450 4800 3300- 7000 4500 82
2 490 7000 5500-10000
6700 83
3 114 4390 3300- 5800
4000 71
4 124 5620 4100- 7200
4700 76
5 149 7090 5200- 9000
6800 73
6 153 5760 3800- 7700 5100 77
ESO Image Processing Group
1. Application Developments
Now that system developments in the
portable version of MIDAS have
stabilized, activities in the area of appli-
cations have resumed.
The echelle package is being up-
graded to minimize the number of para-
meters controlling the reduction se-
quence, and to correct some known de-
ficiencies of the current version. The
new package will also be optimized to
process data from other instruments like
EFOSC and Echelec.
A new package for the reduction of
long-slit spectra is now being tested.
This package will replace the IPCS con-
text of the old MIDAS version.
In the portable version of MIDAS the
implementation of the ROMAFOT
crowded field photometry package has
been completed for DECNMS systems.
Currently an upgrade is in progress to
port the package to UNIX systems. This
upgrade mainly involves the complete
implementation of the MIDAS table file
system and is expected to be released
with the 89 MAY release (see below).
Also in the portable version, the up-
grade of the INVENTORY package was
finalized. The documentation of the
package has been updated accordingly.
A new file system has been im-
plemented in the plot package. In the
new release, MIDAS plot information will
be contained in only one plotfile. This
plotfile is created by the major PLOT
commands and will have the name of
the frame, table, keyword or descriptor
that is plotted. The file extension is
".PLT". Subsequent OVERPLOT com-
mands will append the plot file with the
new plot instructions. The SHOW/PLOT
command shows the user the name of
the current plotfile. MIDAS will allow one
version of a plot file: an old plot file with
the same name as a newly created one
will be deleted. In the SEND/PLOT com-
mand, as the second parameter, one
can specify which plotfile is to be sent to
the graphics device.
2. Data Analysis Workshop
The next Data Analysis Workshop will
be held April 18-20, 1989, in the ESO
headquarters. Its form will change sig-
nificantly in the sense that the main em-
phasis will be placed on astronomical
applications rather than on system re-
lated software. The first day and a half
will be devoted to applications for a
specific area, while the last day will be
used for MIDAS and ST-ECF sessions.
1000 900
will be used at the NTT for tests. Some
scientific work will also be possible, al-
though, because the instrument para-
meters were optimized for the 2.2-m,
EFOSC2 tends to oversample stellar im-
ages. On the NTT the scale at the detec-
tor is 115 which yields a
field of 2:9 x 2:9 with the Thomson
1024 x 1024 chip. With grism # 3 the
slit-limited resolution is 13 Aat 4400 A
with a 1" slil. .
After EMMI has been installed and
tested on the NTT, EFOSC2 will be
moved to the 2.2-m where it will be-
come generally available in the course
of 1990.
::; 40
TABLE 1: Comparison of EFOSC and EFOSC2 optical data.
TABLE 2: Current EFOSC2 grisms.
300 400 500 600 700
WAVElENGTH (nm)
Figure 1: Measured transmission of the optics of EFOSC and EFOSC2.
number of lens groups from 6 to 5, a
careful selection of glass melts and opti-
Cal cements used for the production of
the optics and by shifting the ret/ection
minimum of the single-Iayer MgF2 anti-
rat/ection coating more to the UV. The
Polychromatic image quality is excellent
eVerywhere in the 25 x 25 mm field; 80
cent of the light is concentrated in a
clrcle with a diameter of 20 (0':4).
An initial set of grisms has also been
completed; their properties are
summarized in Table 2. Grisms with
higher dispersion and probably also an
echelle will be added when the final
detector format is known.
In the second half of 1989, EFOSC2
This year the special topic will be analy-
sis of two dimensional direct images
including stellarlsurface photometry,
search for objects and classification.
There will be the possibility of present-
ing short papers during the workshop
and we encourage you to contact us if
you wish to give a contribution or just
participate. We expect that proceedings
of the scientific sessions will be pub-
Iished. Our aim is to create a forum for
discussions of different methods and
algorithms used in image processing.
3. Portable MIDAS
The first official release of the portable
MIDAS, 88NOV, was made with some
delay due to verifications of the VMS
installation procedure. This first version
does not yet include all applications and
especially only supported Gould-De-
Anza . IP 8500 and X-window version
10.4 display systems. It is expected that
these deficiencies will be resolved in the
89 MAY release which will contain the
basic display software for X-window
version 11 being adopted as the stan-
dard for MIDAS. From the 89MAY re-
lease the portable MIDAS will be the
only official version of MIDAS for both
UNIX and DECNMS systems.
4. Access to Astronomical
Catalogues and Databases
A new version (2.2) of STARCAT is
now available; it is accessible from the
ESO computer or remotely through net-
works. Here are some new features of
this version:
• astronomical catalogues can be
queried by a target radius and posi-
tion. The position may be specified in
many coordinate systems (equatorial
at any equinox, galactic, supergalac-
tic, ecliptic);
• J 2000 coordinates are computed
and Iisted for every catalogue;
• the result of any query can be stored
as aMIDAS table, or as a plain ASC11
• remote connections now include
IUE-Vilspa, SIMBAD, and EXOSAT
(ESTEC) but are only available for
local ESO users.
About 30 astronomical catalogues are
available on-line, with complete on-line
documentation. Among the most re-
cently incorporated ones are for exam-
pie the new version (1988) of Abell's
catalogue of clusters of galaxies, and
the most recent version of the catalogue
of White Dwarfs (McCook and Sion,
The same STARCAT interface will be
used for the future catalogue of the ESO
5. Measuring Machine Facility
The central computer of the Measur-
ing Machine Facility is being replaced by
a Stellar GS-1 000 system. The decision
was made after extensive MIDAS
benchmarks giving it the best price/per-
formance. The system will be able to
analyse the scans of full Schmidt plates
which the upgraded OPTRONICS
machine is expected to perform later
this year. The Stellar GS-1000 system
runs a UNIX-like operating system and
has an X-window system version 11 for
display. It will run MIDAS for reductions
of measuring machine data and be con-
nected to the central computers through
Ethernet using TCP/IP protocols.
6. AIPS-MIDAS Agreement
An increasing number of astronomers
are using observations in several
wavelength regions (e.g., optical, in-
frared, and radio) in their research. In
general, different data reduction soft-
ware packages are used for the different
wavelength regions. Also different soft-
ware packages have different capabili-
ties, strengths, and weaknesses. It is
therefore important to ease the transfer
of data between image processing sys-
tems. As a first step in this direction,
AlPS and MIDAS have agreed to write
FITS files on disk with identicial specifi-
cations. This will enable users of these
systems to exchange data files much
faster via disk instead of passing
through a magnetic tape. AlPS already
conforms to the agreement while
MIDAS will implement it as of the
89 MAY release.
The agreement specifies that FITS
disk files have arecord size of 2880
bytes, which is the standard FITS logical
record length. There shall be no "extra"
bytes in arecord, such as those used to
specify variable lengths on some sys-
tems. In this way, FITS disk files may be
passed between different operating sys-
tems through networks with no ambig-
uity. The use of a 2880-byte record im-
plies that reading programmes are not
required to reblock the data into logical
records, although packages such aS
AlPS and MIDAS can be expected to
have that capability in future.
7. MIDAS Hot-Une Service
The following MIDAS support services
can be used to obtain help quickly when
problems arise:
• Tlx.: 52828222 eso d, attn.: MIDAS
• Fax.: +49-89-3202362, attn.: MIDAS
• Tel.: +49-89-32006-456
Users are also invited to send us any
suggestions or comments. Although we
do provide a telephone service we asK
users to use it only in urgent cases. To
make it easier for us to process the
requests properly we ask you, whel'l
possible, to submit requests in writtel'l
form through either electronic networkS
or telex.
Institutes which would like to use the
MIDAS system should submit aMIDAS
Request Form to the Image Processing
Group. This form can be obtained
through the HOT-UNE service.
Ethernet at ESO Headquarters
D. CHIITIM, Benney Electronics, ESO
Ethernet is a cable used with associ-
ated software packages for connecting
computer equipment throughout a
building. It has the advantage that
equipment can be connected along its
length at almost any point. Due to the
high data transmission speeds involved,
ethernet is most suitable for computer-
to-computer communications although
terminals can be connected onto it if
Since the initial installation of ethernet
at the ESO Headquarters two years ago,
the system has grown steadily. There
are now only a few areas of the building
which are not close to the network. The
system consists of both "Thick Ether-
net", which uses an expensive co-a
cable which is almost immune to electri-
cal interference and a thinner cable
which is more susceptible to interfe
ence but cheaper (hence it's nickname
- The main ethernet cable can be five
Computer Room
hundred metres long but may be ex-
tended by using a "repeater".
f - Each device may be no more than
Ifty metres from the cable.
- Each device or "node" on the net-
work has a unique address. When data
IS tranSmitted it is preceded by the ad-
to Which is has to be sent. Each
evice checks the address and will take
only if the data is for that particu-
ar device.
1- The small black rectangles shown
the length of the cables are
. dlum Attachment Units (MAU). These
Plerce the outher shield of the cable and
Contact with the inner conductor.
the case of "cheapernet" this is done
Ith Plugs and sockets).
- Equipment can be connected di-
rectly to a MAU or through a "fan out"
The latter can either be a DEC
ELNI box or an "Isolan" equivalent·
both, however, perform the function of
splitting the connection eight ways.
- The terminal servers provide eight
serial (RS 232) interfaces directly onto
ethernet. One supports the TCP/IP pro-
tocol, the other supports DECNET.
- The "ethernet repeater" serves
mainly as a connection between two
ethernet cable segments whose ends
are inaccessible Networks connected
by a repeater and can be considered to
be a single unit. It can also pravide an
interface between a thick ethernet and a
- The "ethernet bridge" is similar to a
repeater, but it can translate one pra-
tocol into another. It is also used to
separate a network carrying relatively
little traffic fram a heavily used one.
- A thin lan hub gives the possibility
of connecting a thick ethernet cable to
up to four thin ethernet cables.
- 80th ends of every cable must be
terminated with 50 ohms.
4th Floor
I:.;tltemet nt ESO
D. Ch1ttim
"'anuaz:y 1989
Ground Floor
- Two cables to the Vaxes are re-
quired at the moment, one is used for
the interface supporting the TCP/IP pro-
tocol, the other is used for Decnet com-
munications. An increasing number of
devices on the market now support
TCP/IP and only these will be purchased
in future.
- The Technical Division (TD)
"cheapernet" has many devices con-
nected to it including HP computers,
test devices and VME chassis.
Changes to the configuration are
made quite frequently, and it is possible
that this diagram will be slightly in-
correct by the time it is published, but if
sufficient interest is shown I will praduce
updated diagrams from time to time.
My thanks go to Preben Grasb01 and
Charlie Ounnas for their constructive
E In the caption to the reproduction
Narth at Night (see The Messenger
o. 54, p. 15) there was unlortunately no
mention that the reproduction was made
Irom a wall chart 01 Hansen Planetarium.
Earth at Night (23" x 35") is available
Hansen PUblications, 1098 South
West, Salt Lake City, Utah 84101
AX-801-538-2249). Price $ 6.00 +
$ .50 Shipping. Wholesale pricing avail-
a le on request.
The Proceedings 01 the ESO Conlerence on
Very Large Telescopes and their Instrumentation
held lrom 21 1024 March 1988 in Garching, were
published in laie 1988. The Proceedings, which
conlain a total 01 1334 pages, are divided inlo Iwo
volumes and are sold al a price 01 DM 95.-. This
price includes packing and postage (surface mail)
and has 10 be prepaid.
Paymenls have 10 be made 10 Ihe ESO bank
accounl 2102002 wilh Commerzbank München
or by cheque, addressed to Ihe attention 01
Financial Services, Karl-Schwarzschild-Slr. 2
0-8046 Garching bei München.
Please do not lorgel 10 indicale your luli ad-
dress and Ihe lilie 01 Ihe Proceedings.
A Slice of Swiss Cheese Made Out of Steel ...
is suggestive of the main supporting structure of the EMMI speclrograph/imager for the NTT
shown in these two pholographs laken al Iwo phases of the manufacturing al Ihe Oe Pretto-
Escher Wyss factory in Schio, /laly. The lett-hand pielure, laken before welding of Ihe covering
plate, shows Ihe inner slruclure with Ihe complex net of ribs 10 increase Ihe rigidily. The long
poles were inserted temporarily to align the apertures. The right-hand piclure shows Ihe
welded, cleaned piece mounled on a measuring machine for the check of Ihe dimensions. The
openings correspond 10 Ihe main oplical componenls 10 be inserted: collimators, filler and
grism wheels, folding mirrors, eie.
The dimensions of EMMI's main slruclure are 240 cm in lenglh, 160 in heighl and 28 only in
depth; Ihe weighl is aboul 800 kg.
The slruclure will arrive al ESO in February. Mounling of the different funclions (now being
inlegrated and lesled in Garehing) on Ihe main body and lesting of the overall instrument will
lake place wilhin Ihe summer of Ihis year. S. D'ODORICO, ESO
H. Schwarz: The Southern Crab . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 1
H. van der Laan: Key Programmes on La Silla: First Allocations . . . . . . . . . . . . . . . . . . 2
V. de Lapparent et al.: Profile of a Key Programme: A Redshift Survey of Galaxies with z::S
0.6 Using Multi-Object Spectroscopy . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 5
R. Bender et al.: Profile of a Key Programme: Towards a Physical Classification of Early-
type Galaxies . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 6
J. Surdej et al.: Profile of a Key Programme: Gravitational Lensing 8
H. van der Laan: The New Research Student Programme of the European Southern
Observatory . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 12
B. Marano: The Users Committee (UC) . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 13
Tentative Time-table of Council Sessions and Committee Meetings in 1989 . . . . . . . . . . . . 14
A. Blaauw: ESO's Early History, 1953-1975. 11. Searching for a Site in South Africa . . . . . . 14
L. Zago and F. Rigaut: Wind Turbulence in the Dome of the 3.6-m Telescope . . . . . . . 22
Visiting Astronomers (ApriI1-0ctober 1, 1989) . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .. 24
T. Le Bertre et al.: IRC+ 10216: a Peanut Nebula! . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .. 25
List of ESO Preprints (December 1988- February 1989) 28
Guidelines for Authors of Articles for the ESO Messenger .. . . . . . . . . . . . . . . . . . . . . . . . . 28
W. C. Keel: A Distant View of La Silla . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .. 29
Staf"f Movements . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .. 29
C. Madsen: The Changing Skyline of La Silla 31
B. Reipurth and C. Madsen: Signposts of Low Mass Star Formation in Molecular Clouds. 32
E. Bajaja and E. Hummel: Observation of the 12CO (J =-> 0) Line in NGC 613 with the
R. Bandiera and P. Focardi: High Resolution Ha Spectroscopy of Nova Centauri 1986:
Tracing a, Transient in the Spectral Evolution. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .. 39
G. Mathys et al.: The European Working Group on Chemically Peculiar Stars of the Upper
Main Sequence: The First 10 Years . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .. 41
J. Andersen et al.: AK Scorpii: A New Pre-Main-Sequence Spectroscopic Binary 45
B. M. Shustov: SupersheIls and Galactic Fountains . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .. 47
F. Paresce: Searching for Light Echoes from Circumstellar Dust Shells Around SN 1987A. 49
G. Burki et al.: The Geneva Photometrie Monitoring of SN 1987A . . . . . . . . . . . . . . . . . . .. 51
Is there a Pulsar in Supernova 1987A? . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .. 52
H. Butcher and 1. Schoenmaker: Squeezing the Most from the CES. . . . . . . . . . . . . . . . .. 52
G. Meylan et al.: What is the Mass-to-Light Ratio of the Old Magellanic Globular Cluster
H. Barwig: High-tech Telescope on Top of Mount Wendel stein . . . . . . . . . . . . . . . . . . . . .. 59
E. Giraud: Spatial Resolution Imaging on the Radio Source 3C 255 60
R. Haefner: The Spectacular Binary System PG 1550+131 .. . . . . . . . . . . . . . . . . . . . . . .. 61
G. Avila et al.: Report on the Last Observing Run of Multiobject Spectroscopy: OPTOPUS
is Alive and Kicking 62
H. Dekker: Optics and Grisms of EFOSC 2 . . . . . . . . . • . • . . . . . . • . . . . . . . . . . . . . . . . . 64
ESO Image Processing Group: MIDAS Memo. . . . . . . . . . . . . . . • . . . . . . . . . . . . . . . . . 65
D. Chittim: Ethernet at ESO Headquarters ................•......•............ 66
S. D'Odorico: A Slice of Swiss Cheese Made Out of Steel . . . . • . . . . . . . . . . . . . . . . . . . 68