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A DETERMINATION OF KINEMATICS AND STRUCTURE
OF THE INNER GALACTIC HALO USING
BLUE FIELD STARS AS TRACERS OF
LARGE-SCALE HALO PROPERTIES

Steve Peter Doinidis

A THESIS

Submitted to
Michigan State University
in partial fulfillment of the requirements
for the degree of

MASTER OF SCIENCE

Department of Physics and Astronomy

1991

ABSTRACT

A DETERMINATION OF KINEMATICS AND STRUCTURE
OF THE INNER GALACTIC HALO USING
BLUE FIELD STARS AS TRACERS OF
LARGE-SCALE HALO PROPERTIES

by

Steve Peter Doinidis

The structure and kinematics of the inner galactic halo are explored by utilizing
samples of blue field stars as tracers of large-scale halo properties. Using the cata-
logue of blue field horizontal-branch (FHB) stars of Beers, Preston and Schectman,
it is demonstrated that the distribution of these stars is not uniform, and that struc-
ture exists in their distribution on scales of less than 100 parsecs. The extension
of the HK Survey of Beers, Preston and Schectman to the northern galactic hemi-
sphere is discussed, outlining the methodology of both the survey and the follow-up
observations that are made on the blue halo stars identified therein. A sample of
592 stars for which spectroscopic measurements have been taken, drawn from the
HK Survey in the southern galactic hemisphere, consisting of 353 FHB stars and
239 main sequence A-type stars, are applied to the kinematic models of Frenk and
White to determine the rotational velocity of these field star systems. These stel-
lar types are treated both independently and in combination for this analysis, and
the results are compared with previous analyses which have been made for other

systems in the halo.

This is dedicated to my family,
and to the memory of my grandfathers,
Steve Doinidis and Zachariah Kourous.

ACKNOWLEDGEMENTS

I would like to acknowledge Dr. Timothy Beers,

whose patience and instruction made this work possible.

I would also like to acknowledge Dr. Greg Kenning,
Benita Pena, Melton Towns, and Stephanie Holland
for keeping my spirits buoyed and my sanity intact
over the duration of my studies at Michigan State.

ii

LIST OF TABLES

LIST OF FIGURES

TABLE OF CONTENTS

KEY TO SYMBOLS, ABBREVIATIONS AND NOMENCLATURE

CHAPTER 1

CHAPTER 2

CHAPTER 3

CHAPTER 4

CHAPTER 5

CHAPTER 6

CHAPTER 7

INTRODUCTION

STATISTICAL TESTS FOR SPATIAL
STRUCTURE IN THE HALO

EVIDENCE FOR CLUSTERING OF
FIELD HORIZONTAL-BRAN CH STARS
IN THE GALACTIC HALO

THE EXTENSION OF THE HK SURVEY
TO THE NORTH GALACTIC HEMISPHERE

PHOTOMETRY OF A SAMPLE OF STARS
FROM THE NORTHERN HK SURVEY

A DETERMINATION OF THE
ROTATIONAL VELOCITY OF A SYSTEM
OF FIELD STARS

SUMMARY

LIST OF REFERENCES

iii

iv

vii

14

LIST OF TABLES

TABLE 1: HK Survey Stellar Classifications.

TABLE 2: Line Index Wavelength Bands.

TABLE 3: Spectroscopy and Photometry of FHB and A Stars.

TABLE 4: FW80 Analysis Results.

TABLE 5: FW80 Analysis Results.

iv

29

32

68

92

94

LIST OF FIGURES

FIGURE 1: Unreddened two-colour diagram. The solid curve repre—
sents the main sequence relation, the dashed curve synthetic colours
from Buser and Kurucz (1978) for log g = 4.5 and [Fe/H]: -2.5. The
stellar types represented here are FHB stars (filled circles), candidate
metal-poor halo stars (diamonds) and field stars of intermediate type
(crosses). The lower line represents the reddening. Figure reproduced
from Preston at al. (1990).

FIGURE 2: A plot of H 7 vs. H 6 for hot blue field stars. Figure
reproduced from Beers et a1. (1991).

FIGURE 3: A plot of D0,; vs. RC for a variety of hot stars in the halo.
Region (I) is occupied by white dwarfs, st and sdB stars, region (II)
by st, sdB and high luminosity stars, region (III) by FHB and A
type stars, and region (IV) by white dwarfs. Figure reproduced from
Beers at al. (1991).

FIGURE 4: Theoretical relation for KP vs. unreddened (B - V),
showing loci of [Fe/H]: -0.5 (upper line), [Fe/H]: -1 (middle line)
and [Fe/H]: -2 (lower line).

FIGURE 5: Theoretical relation for (B - V) vs. HP for [Fe/H]=-1.
The solid line represents the relation for log g = 3 and the dashed
line, log 5 = .

FIGURE 6: Distributions in position of FHB stars (left side) and A
type stars (right side) in the Z—X and Y~X planes.

FIGURE 7: Distance above the galactic plane, Z. vs. abundance
for the FHB stars (open circles) and A type stars (filled circles). The
dashed lines represent a typical distance for each of these types near
the magnitude limit of the HK Survey.

FIGURE 8: Definitions of the geometric and physical parameters used
in the FW80 method (Figure reproduced from FWSO).

34

35

36

37

65

66

89

FIGURE 9: Solution for v,,,, 010., and their ratio as a function of
abundance, using the method of FW80 and the data of Table 4.

FIGURE 10: Solution for v,.,,, 0'1”, and their ratio as a function of
abundance, using the method of FW80 and the data of Table 5.

FIGURE 11: A comparison of derived values for v,“ for three different
divisions between the FHB and A type star samples: [Fe/ H] = —0.5
(top), [Fe/H] = —0.75 (middle) and [Fe/H] = -1.00 (bottom).

FIGURE 12: Stripe histograms showing the range in cos 2,12.- for each
of the bins in Table 4.

vi

93

95

97

99

KEY TO SYMBOLS, ABBREVIATIONS AND NOMENCLATURE

GCS - Globular Cluster System.

FW80 - Frenk and White (1980).

T89 -— Thomas (1989).

BPS I - Beers, Preston and Schectman (1985).
BPS II — Beers, Preston and Schectman (1990).
FHB — Field Horizontal-Branch.

FHB I — Beers, Preston and Schectman (1988).
SLC - Sommer-Larsen and Christensen (1987).
PCF - Pairwise Correlation Function.

KPNO - Kitt Peak National Observatory.
CTIO - Cerro Tololo Interamerican Observatory.
KP - Line index based on K feature of Call.

HP - Line index based on H 7 and H6.

D03 — Width of H6 feature at 20% below the level of the local continuum.
RC - Depth of H6 feature relative to continuum.

vii

CHAPTER 1: INTRODUCTION

The study of the galactic halo offers clues to the formation and evolution of
the galaxy. This is due to the unavoidable chemical and dynamical mixing that
has occurred in the disk since its formation, virtually eliminating the record of
the initial collapse and early stages of the protogalaxy that formed our present
system. Because the halo is much less dense than the disk, there has been little
opportunity for the individual components of the halo system to interact with each
other, and it is uncertain even if the halo has dynamically relaxed since the collapse
and formation of the disk system. If this is indeed the case, then objects in the
halo can be utilized as tracers of the formation of the galaxy. The halo can also
provide information on the kinematics and dynamics of the galaxy to large distances
from the galactic center, free from the obscuration of the gas and dust in the disk.
Objects in the halo can be used to empirically establish the form of the galactic
potential to large galactocentric distances, which allows the testing of models for
the kinematical behaviour of the galaxy and work towards refinements of these
models. These observations can also provide information on more general problems
in galactic astronomy, such as the large-scale structure of spirals and the dark matter

hypothesis.

2

The globular cluster system (hereafter GCS) was the first and most obvious
system of halo objects to be studied. Kinematic studies, such as those of Kinman
(1959), Clube and Watson (1979) and Henk and White (1980, hereafter FW80) have
formed the basis of our understanding of the processes which govern the nature of the
halo. Studies on the spatial distribution of the GCS (e.g., Woltjer, 1975, Thomas,
1989, hereafter T89) suggest that the form of the galactic potential is spherical,
but it is inherently problematic to apply the GCS to these ends, as the number of
clusters which can be observed number in the hundreds while the volume occupied
by the system is on order 105 — 106 kpc3. This limits the conclusions that can
be drawn from these studies to areas of large scales and meager detail. Another
limit of the globular cluster system is the large uncertainty in distance estimates
for many clusters, which have distance modulii accurate only to about 0.5 mag, a

figure unlikely to improve in the near term.

It would be far more advantageous to employ halo stars in determinations of
halo structure and kinematics, as they are far more numerous than the globulars.
They are also distributed more uniformly throughout the halo, giving much more
predictive power to the observations that can be made on them. This advantage
is offset, to a degree, by their faintness, which makes it dificult to unambiguously
identify field stars in large-scale surveys. The first attempts to collect samples of
halo stars involved the identification of high velocity stars, as well as surveys for

faint blue stars in the direction of the galactic poles.

The high velocity stars are so—called because they are close to the solar neigh-

3
bourhood, yet have heliocentric radial velocities that are in excess of the circular
velocity that a star would have at the sun’s galactocentric distance, the inference
being that their orbits are much more energetic than that of the sun, making them
members of the halo population. Because the criterion used to select these stars
introduces an obvious kinematic bias to any samples collected in this manner, their

usefulness in applications to dynamic modelling is limited.

Surveys for faint blue stars have relied on two basic methods for collecting
samples of stars. The first is simply to look for stars with blue colours and faint
limiting magnitudes. Stars are identified either through wide-field photographic
plates taken through broad-band filters, or by individual photometric measurements
on selected stars. The colour of the objects is, however, a weak criterion from which
to select field stars, and there are problems with contamination of these samples
from subluminous stars in the stellar neighbourhood. The range of stellar types
which can be found is also limited to the hottest of the field populations, which are
also the least numerous. A similar method utilizes the colour differences of stars,
(B — V) and (U - B), which can be used to place a star on a two-colour diagram
(see Figure 1). This entails either measurements of photographic magnitudes of
these stars on plates taken with broad-band filters, which are accurate only to :i:0.2
mag, or by individual photometric measurements, which are inefficient and time

consuming.

Over the last ten years, two survey methods have been developed to more easily

identify large numbers of halo stars by making use of the large areal coverage that

 

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‘ L L 1 I L 1 1 L L L L l L L L L L 1 l I L I L
'-2 o 2 .4 .s .a 1.0
(B-V).

 

Figure 1. Unreddened two-colour diagram. The solid curve represents the main
sequence relation, the dashed curve synthetic colours from Buser and Kurucz (1978)
for log g = 4.5 and [Fe/H]: -2.5. The stellar types represented here are FHB
stars (filled circles), candidate metal-poor halo stars (diamonds) and field stars
of intermediate type (crosses). The lower line represents the reddening. Figure
reproduced from Preston ct a1. (1990).

5
is possible with Schmidt telescopes. The first method is to search for stars with
large ultraviolet excesses, and has been employed successfully in the the Palomar
Green Survey (Green 1980), and in the Kiso Survey (Noguchi, Maehara and Kondo
1980). The methodology for these surveys is essentially the same as earlier surveys
that sought to identify blue stars at high galactic latitudes, but better techniques
have since been developed for colorimetric classification of these objects into the
various stellar types which populate the halo (Kilkenny 1987). Spectroscopic and
photometric follow-up observations are nevertheless required in order to confirm
the candidate stars that are identified in these surveys, and so they are by nature

long-term projects.

The second method is to search for candidate stars by analyzing their spectra
on objective-prism plates taken on these telescopes. Earlier surveys, such as those
of Slettebak and Stock (1959) and Slettebak and Brundage (1971) concentrated on
identifying early-type stars in the vicinity of the galactic poles over the magnitude
range 6 < B < 14 — 15. Some difficulty is faced at the faint magnitudes in these
surveys, owing to overlap between adjacent spectra. At the magnitude limits, this
effect is very pronounced, making it dificult to accurately identify stellar types
which are of interest. The HK Survey of Beers, Preston- and Schectman (1985,
1991; hereafter BPS I and BPS II) utilizes an objective-prism coupled with an
interference filter to identify blue halo stars, particularly halo stars of extremely
low metal abundance. The filter has a narrow (a: 150 A) bandpass centered in

the region of the H and K lines of Call, a wavelength range with features that are

6

suficient for the identification of these stars with a high degree of success.

By using such a narrow bandpass for the plates, the resultant spectra are them-
selves smaller (although they are widened for clarity), so that crowding of spectra
is not as great a concern as for other methods, allowing for a fainter limiting mag-
nitude, covering the range 11 < B < 16. This work will be concerned primarily
with those stars in the survey which were classed as type A and AB, for which
the great majority (Z 85%) are expected to be blue field horizontal-branch (FHB)
stars. Among the other types expected are A type stars, hot subdwarfs (st and
sdB) and degenerates. These stars alone are numerous enough to allow for a com-
prehensive determination of all kinematical and dynamical parameters in the halo.
Another major advantage of this survey method is the efficient use that it makes
of observational facilities. This is due partly to the large areal coverage of Schmidt
plates, but mostly due to the high percentage of the candidate stars which are
identified correctly on these plates. This ensures that the majority of the follow-up

spectroscopic and photometric observations can be put to use.

A catalogue of FHB candidate stars identified through the HK Survey in the
southern galactic hemisphere has been presented by Beers, Preston and Schectman
(1988, hereafter FHB I), containing some 4400 candidate stars. The data included
for each star consists of coordinates accurate to 2 arcseconds, and rough magnitude
estimates based on the brightness of the spectra, accurate to about i0.5 mag.
Although the uncertainty in magnitude is roughly equal to that for the globulars,

the prospects for improving these are limited only to the availability of observational

7
facilities for photometric measurements of these stars. Distances for these stars are
easily determined, owing to the well-defined intrinsic brightness of the horizont al-
branch (Mv = 0.6 :i: 0.2 mag). Chapter 2 outlines the use of an estimator, the two-
point correlation function, which can be used to probe the distribution of the FHB
stars in the catalogue for deviations from that expected for a random distribution
of stars. In Chapter 3, it is demonstrated that the FHB I sample indicates that
there is structure in the distribution of the halo stars on scales of less than 100
parsecs, suggesting that models for the galaxy which assume dynamical relaxation

in the halo may not be truly reflective of the actual conditions existent in the halo.

In Chapter 4, the extension of the HK Survey into the northern galactic hemi-
sphere is discussed, as well as the determination of the stellar types of candidates
based on follow-up spectroscopic and photometric observations is detailed. Because
the scanning of the northern survey plates is carried out by Beers as opposed to
Preston, who carried out the scanning in the south, it is important to establish that
the methodology employed by both is consistent. It has therefore been important to
initiate the follow-up observations of the northern hemisphere candidates in order to
be able to address any possible differences. In Chapter 5, photometry for a sample
of stars drawn from the northern survey is presented, and the issue of uniformity

between the northern and southern samples is discussed.

In Chapter 6, a sample of 585 stars from the southern HK Survey for which
spectroscopic measurements have been made are applied to the kinematic models

of FW80 to determine the rotational velocity of the inner halo. The sample consists

8
of 375 FHB stars and 210 stars with normal A type spectra. The stars are applied
both as separate samples and a combined sample, and the results are compared

against similar determinations made for other halo systems.

CHAPTER 2: STATISTICAL TESTS FOR SPATIAL STRUCTURE

IN THE HALO

The spatial structure of the halo has for some time been a topic of great interest
and a property of the galaxy which is poorly constrained. The interest in this facet of
the halo comes from the link that the spatial distribution of objects could have with
the overall mass distribution of the galaxy, so that knowledge of this distribution
would in turn provide us with the form of the galactic potential. The result which
is obtained using the GCS is consistently that of a spherical distribution, but over
the large scales which are necessitated by using this sample. It is unclear as to
whether the GCS can adequately constrain the geometry of the mass distribution
on scales smaller than a few kpc (Woltjer, 1975). This is due in part to scale, but
also to differences between individual globulars (e. g., metal-poor halo clusters vs.
metal-rich disk clusters) which make it dificult to draw any direct conclusions from

the entire sample.

Studies have also been made using field star samples, and in general the re-
sults have depended on whether kinematics or star counts were used in making the

determination. High-velocity, metal-poor subdwarfs in the solar neighbourhood

10
exhibit a velocity dispersion tensor 0,, : 099 : a" z 2 : 1 : 1, suggesting that, if
these stars are tracers of the mass distribution in the halo, the potential is flattened,
at least to distances of a few kpc from the sun (Gilmore at 0.1., 1989), with an axial
ratio of c/a z 0.6. Star count studies are more consistent with a distribution that
is close to spherical. By comparing star counts in two fields at different latitudes
along the l = 90°, 270° plane selected by K00 and Kron (1982), Bahcall and Soneira
(1984) derive an axial ratio of c/a z 0.80:8:32. In similar fields (but at fainter
magnitudes), Koo et al.(1986) derive a value of c/a z 1.3, but this figure is not as

well constrained.

In general, much of this work rests on the assumption that these objects are
tracers of the halo mass, when in fact this is not clearly evident on small scales.
Much, if not most of the halo mass is in the form of dark matter, and so care
must be taken to utilize samples that cover a wide range of positions in the halo in
order to properly account for this. The announcement of the discovery of a group of
FHB stars which were clustered together and which had similar velocities (Sommer-
Larsen and Christensen, 1987, hereafter SLC) raised the possibility that there could
be gravitationally bound structures in the halo on small scales, independent of the
GCS. This group exhibited a spread in distance modulus of 0DM = 0.16 and a
velocity dispersion of a” 5 20km 8“, so that if the FHB stars are tracing a larger
group, then it would have a characteristic size and mass of a globular cluster with

an extremely large (a: 500 : 1) mass-to-light ratio. Recent work on the SLC group

(Sommer-Larsen and Christensen, 1989) has shown that one of the five stars in the

11
group is a main sequence gravity A star, and that a more accurate determination
of the velocity dispersion gives 0,, = 53 :I: 22km 8-1, which severely weakens the

possibility that the group is bound.

The original finding did, however, motivate a search through the data of Pier
(1982, 1983) revealing a similar grouping of FHB stars (Beers and Doinidis, 1988),
and this raised the possibility that there could be additional such groups present
in the halo. As the data of Pier was drawn from candidate A and AB type stars
identified in the HK Survey, it is possible that additional groups of this kind could
be found in the FHB I catalogue. This catalogue contains positional data and rough
estimates of apparent brightness (:l:0.5 mag) for the FHB candidate stars contained
within it, so it is possible to test, albeit tentatively, for groupings of these stars in
space on scales similar to that found in the SLC and Pier groups. Confirmation of
these groups as dynamically bound entities would require accurate photometry and

radial velocity measurements.

Groups such as these, if they are prevalent, would seem to indicate that the
distribution of halo stars is not uniform at some scale, and this could be determined
given a large enough sample to test this hypothesis. A number of statistical method-
ologies are available to explore this possibility, and the choice of estimator is driven
mainly by the nature of the sample to be studied. The FHB I catalogue consists of
candidate FHB stars identified on individual Schmidt plates, which do not form a
contiguous area on the sky (see Figure 1 of FHB I). The sample, therefore, cannot

be treated as a whole. It is possible, however, to treat the plates individually and

12
combine the results derived in this manner for the individual plates. An estima-
tor which can be applied in this way is the pairwise correlation function (hereafter

PCF), which has been previously used to explore clustering among galaxies.

A Monte Carlo PCF estimator, w(0), can be defined by comparing the number
of pairs of objects, N,(0), which have angular separations between 0 — A0/2 and
9 + [39/ 2 found in the FHB I fields to Nr(9), the number expected from a random

distribution, such that

Np“)

w(0)=m-1i (1)

Because this estimator requires no separate calculation for edge effects which are
introduced by setting an upper limit to 0, it can be used with sample fields of any
expedient size. This allows its appplication to the FHB I sample as a collection of
individual plates. If one assumes a spherically symmetric, random distribution for
these stars, then the PCF derived for each of the plates should be identically zero.
Any statistically significant deviation from zerois then evidence that some structure
exists. The test for this hypothesis consisted of a comparison of the plates against
the expectation for a random distribution, which was determined by creating 100
random realisations of each of the survey plates. It was found that the PCF for
the catalogue differed from that for the random realisations to a significant degree,
implying that the halo contains structure that deviates from that expected for a

random distribution. A summary of this work is given in the following chapter.

13

This alone would not imply that other clusters of FHB stars necessarily exist,
or that these clusters can account for the observed PCF. A further test using a
single-linkage clustering algorithm was applied to the FHB I data. This algorithm
searches around a specific star for other stars that satisfy the linkage criterion,
namely, that the other stars are within a specified distance. This criterion was set
to the typical separation between stars in the SLC and Pier groups, and by this
method, 23 candidate groups were found, compared to an average of 12 for the
randomly generated catalogues. It was also determined that the candidate groups
do not fully account for the observed strength of the PCF. Finally, it should also
be noted that SLC explore the possibility of clustering in velocity space for their
fields in the direction of the north and south galactic poles, finding no groups of
stars with velocity dispersions less than 20 km 8'“. Because of the small numbers
of FHB stars expected in such high latitude fields, it is uncertain as to whether this

demonstrates a general trend for the halo.

CHAPTER 3: EVIDENCE FOR CLUSTERING OF FIELD

HORIZONTAL-BRANCH STARS IN THE GALACTIC HALO

Published in The Astrophysigl Jourrg (Letters),

Volume 340, Pages L57 to L60.

14

15
EVIDENCE FOR CLUSTERING OF FIELD HORIZONTAL-BRAN CH STARS
IN THE GALACTIC HALO
Steve P. Doinidis and Timothy C. Beers
Department of Physics and Astronomy, Michigan State University

ABSTRACT

We report on an investigation of the clusteringifield horizont al—
branch stars in the Galactic halo. We examine this question using a cat-
alog of over 4400 candidate field horizontal-branch stars distributed over
roughly 2300 square degrees of sky, primarily in the southern Galactic
hemisphere. A two—point correlation analysis indicates that the catalog
contains an excess of stellar pairs with angular separations 0 S 10 are
minutes; at the approximate distances of these stars (5-8 kpc from the
sun) the corresponding linear separations are r S 10 pc. We comment on
the possible connection with loose groups of halo field horizontal-branch
stars noted previously by Sommer-Larsen and Christensen, and Beers and
Doinidis.

I. INTRODUCTION

This Letter reports on an initial investigation of the clustering properties of
field horizontal-branch (hereafter FHB) stars in the Galactic halo. This study was
motivated, in part, by the report of an apparent physical group of FHB stars dis-
covered in a limited sample of blue objects on a single wide-field Schmidt plate _
(Sommer—Larsen and Christensen 1987; hereafter SLC). The proposed group con-
sists of five stars with a spread in apparent magnitude av = 0.16 mag, and a radial
velocity dispersion a, S 20 km s". The spatial extent of this group (roughly 100
pc at a distance of 4 kpc from the sun) prompted SLC to speculate that these stars
are part of a high mass-to-light object with dimensions of a globular cluster, or
remnants of a recently disrupted system. In a sub-sample of some 200 FHB stars
with complete photometric and radial velocity measurements, Beers and Doinidis

(1988) identify a similar group with at least five members, which suggested that
other such groups might be found.

FHB stars serve as valuable probes of the Galactic halo. They are intrinsi-
cally luminous (M v z 0.8), homogeneously distributed, and easily identified, either

16

spectroscopically (Pier 1983) or photometrically (Schechter 1988). They are also
numerous. Preliminary investigations of the number density of FHB stars in the
region of the Galactic bulge indicate that they outnumber the RR Lyrae variables
in the same field by roughly a factor of 5 to 10 (Preston 1988).

Large samples of FHB stars enable investigations of the kinematics and dynam-
ics of the halo on much smaller scales than heretofore possible. Previous searches
for faint blue objects, primarily in the direction of the Galacticip'oles, have identified
several hundred early-type stars, some of which might be expected to be halo FHB
stars (Chavira 1958; Ram and Luyten 1962; Philip and Sanduleak 1968; Slettebak
and Brundage 1971). The catalog of Beers, Preston and Schectman (1988; hereafter
FHB I) contains some 4,400 FHB'candidates selected from an objective—prism sur-
vey of over 2300 square degrees in the south Galactic hemisphere. Spectroscopy of a
limited sample of these candidates indicates that over 85% are bona-fide members
of the halo horizontal-branch population (Pier 1983). Details of the ob jective—prism

survey itself are given in FHB I.

As discussed in FHB I, apparent magnitudes estimated from the density of
individual objectiveprism spectra were assigned to each candidate, with an error
on order 0.5 mag. The distance to each candidate is derived from its estimated
apparent magnitude, assuming a horizontal-branch luminosity of M v = 0.8 . Color
information is not available for most candidates, and no attempt is made to correct
the inferred distance for the color dependence of the horizontal-branch luminosity.
It is apparent from the available photometry (Pier 1983; Preston 1988) that the
great majority of candidate stars in FHB I lie in the color range ~—O.1 S B -V S 0.2 ,
where no appreciable color correction is expected. Small shifts in horizontal—branch
luminosity might be expected if the helium abundance of the FHB population varies
over a wide range (Sweigart 1985), but this effect is small compared to the random
error in estimating apparent magnitudes. Because the plates are primarily located
at high Galactic latitudes, interstellar reddening is not expected to be appreciable.

17

II. THE TWO—POINT CORRELATION ANALYSIS

We restrict our attention to the candidates occupying the three most populous
brightness classifications (FHB I classes m, mf, and f); these classes correspond
to median distances of 5, 6, and 7.5 kpc from the sun, respectively. It should be
emphasized that the 92 plates available to date do not represent a contiguous sample;
each plate is 5 degrees in diameter with little or no overl;.p\even in well-surveyed
portions of the sky. We simulate the FHB I catalog by replacing the coordinates
of each candidate on a given plate with draws from a random distribution. This
procedure retains the original brightness distributions of the candidates on each
plate. One hundred random realizations of the FHB I catalog were created.

Significant small-scale clustering should be revealed by power in a two—point
correlation plot of angular separations. Following the method of Hewett (1982) we

obtain:

«ohm-1

Nr(9) . (1)

where N,(0) is the number of pairs in the FHB I sample with separations in the
range 0d: A0 and N,.(0) is the corresponding number of pairs in each of the random
catalogs. This procedure eliminates the need for explicit correction of edge effects,
which might be expected to be severe due to the low areal density of stars on each
plate. The faintest FHB candidates we consider are at least 1 magnitude above the
plate limit, thus small variations in plate sensitivity should not affect the sample.
We choose A0 = 5 arc minutes, and count pairs to separations of 90 are minutes,

without distinguishing between stars of different brightness classes.

A matter of concern in the correlation analysis was the possibility of bias due
to the inclusion of FHB stars in the outskirts of globular clusters. The presence
of a few such interlopers would increase the number of close pairs in our sample,
and completely dominate the signal. We have checked the list of Harris and Racine
(1979) for globulars within 3 degrees of the centers of the catalog plates. We found

18

only four plates on which there was the possibility of contamination of our sample.
In only one of the four cases (CS 22890) was the nearby globular cluster at a
distance where its members might overlap in apparent brightness with stars in
FHB I. Horizontal-branch stars in M5 have apparent magnitudes corresponding to
FHB I brightness class f. We exclude the 14 likely members of M5 before making
our calculations.

\K.
The final correlation we obtain is the mean of the correlation functions calcu-

lated by comparing the FHB I sample with each of 100 random catalogs. Errors
bars are obtained from the one—sigma scatter about this mean. Making use of the
approximate distances assigned to each discrete brightness class, we calculate the
corresponding linear correlation function:

50-) = %% — 1. (2)
We use bins of width 10 pc, to separations of 100 pc. Only stars of the same
brightness class were counted as possible pairs in this case. Figure 1a shows the
mean angular correlation function for stars in the FHB I catalog. The solid lines
indicate the onrsigma variation about a purely random distribution. In the range
0 S, 9 _<_ 10 arcminutes there is a weak, but significant, correlation (amplitude
z 0.40 d: 0.03). In this bin, 488 unique pairs in the FHB I catalog were identified
compared to an expected 348 pairs from the simulations. A statistically significant

correlation exists out to separations on order 60 arc-minutes.

The linear correlation is shown in Figure 1b. Again, the signal in the first
bin is weak but significant (amplitude z 0.40 :I: 0.12), resulting from 186 pairs in
FHB I compared to an expected 133 pairs. No significant correlation is observed at
larger separation. The large scatter in the linear correlation is expected due to the

coarseness of the brightness estimates.

We have considered the possibility that the observed correlation is due, in part,
to potential selection biases in the original classification of the FHB candidates. The

19

spectroscopic signature of the FHB stars on the objective—prism plates is unmis-
takable - a weak or absent CaII K line accompanied by a strong He absorption
feature and a blue continuum. The only doubts concerning classification are for
the very faintest stars (class vf), which are excluded from our correlation analysis.
It is possible that the brightness classes which are assigned to each candidate star
are influenced by the proximity of one star to another, which might potentially
bias the linear correlation calculation, but of course shouthdt' affect“ the angular

correlation.

III. A SEARCH FOR CANDIDATE GROUPS

The observation _ of a significant two-point correlation in the positions of the
FHB stars is a necessary, but not suficient condition for the existence of loose groups
such as those identified by SLC and Beers and Doinidis (1988). For example, the
correlation power could be dominated by isolated pairs or triplets rather than groups

of five or more objects.

Given the rough magnitude estimates presently available for the stars in FHB I
'we are hesitant to present a definitive list of FHB groups. A strict evaluation of the
significance of a group-finding procedure would require a model for the subjective
process of assigning discrete brightness estimates for each star, which would be
rather ad—hoc. We only comment that an application of a single—linkage clustering
algorithm (based on the code of Huchra and Geller 1982) on the FHB I catalog
detected 23 groups of 5 or more stars in identical brightness classes; the same
algorithm identified, on average, 12 such groups in the simulated catalogs.

One group of FHB candidates identified by the group-finding algorithm, dis-
cussed by Beers and Doinidis (1988), is very similar to the SLC group. This group
is identified on a plate at relatively low galactic latitude (b = -15°) in a direc-
tion toward the Galactic center. The list of potential members, along with avail-
able velocity and photometric information, is provided in Table 1. Column 1 is
the star number from FHB I. Columns 2 and 3 list the 1950 equatorial coordi-
nates. Columns 4—6 give the apparent magnitude and colors (if available) from Pier

20

(1982). The distance moduli for candidate members are calculated using the pa-
rameterization of horizontal—branch absolute magnitude from Sommer—Larsen and
Christensen (1986), and are given in column 7. Heliocentric radial velocities (Pier
1983) are listed in column 8. Column 9 lists the equivalent width of the Call K
feature (also from Pier 1983).

Two stars are rejected from consideration on the basi\sof their large positive
velocities. The group covers a total area of roughly 1 square L degree of sky. The
mean pairwise separation of candidate members is 33 arc—minutes, corresponding to
roughly 45 pc at the median group distance of about 5 kpc from the sun. The pro-
posed members with measured photometry have a dispersion in apparent magnitude
av = 0.14 mag. The corresponding spread in distance moduli is a'DM = 0.12 mag.
The radial velocity dispersion for the five stars judged to be members is a, = 21
km 8.1. . We caution that the radial velocities have external errors no better than
10 km s"; the SLC measurements are even less accurate. It is therefore possible

that the actual dispersions of the groups are substantially smaller than the reported
20 km s". More accurate velocities are clearly required.

IV. DISCUSSION

The existence of structure in the halo FHB population, if confirmed by fur-
ther analysis, is of great importance for constraining models of the formation and
evolution of the Galaxy. For example, in the presence of a population of massive
black holes, which have been suggested as a major component of the dark coronac
of galaxies (Lacey and Ostriker 1985), loose clusters with masses and radii typical
of the halo FHB groups would be rapidly disrupted (Wielen 1987). On the other
extreme, it is even possible that the apparent high mass—to—light ratios of the FHB
groups indicate that they, themselves, are harboring the proposed black holes.

The FHB groups are not the first suggested clusters of halo stars other than
the globulars. Proper motion and radial velocity surveys of stars in the solar neigh-
borhood suggest that at least some “moving groups” of stars exist (Eggen 1987;

21

Ratnatunga 1988). One such group, Groombridge 1830, has a measured space ve-
locity large enough to preclude its membership in the disk population (Eggen and
Sandage 1959).

We are still ignorant of many basic properties of FHB groups. It is not even
known, for example, whether these groups have corresponding populations of other
kinds of stars, such as red giants and main sequence stars. Freeman (1988) is in-
vestigating this question for the SLC group; a result shoul'dgghortly, forthcoming.
We plan on obtaining photometry and radial velocities for all pairs and members of
the proposed groups in the FHB I catalog. With this information in hand, a more

definitive evaluation of structure in the halo FHB population can be made.

90

~-

REFEREN CES

Beers, T.C., and Doinidis, S.P. 1988. Bull.A.A.S., 19, 1030.

Beers, T.C., Preston, G.W., and Schectman, SA. 1988, Ap.J.Suppl. 67,
461 (FHB I).

Chavira, E. (1958), 301. Obs. Tonantzintla y Tacubaya, 2, No. 17, 15.

Eggen,O.J. 1987 in Proc. NATO Advanced Study Institute on The
Galaxy, ed. G. Gilmore and B. Carswell (DordrechtzReidel),
p.211.

Eggen, O.J., and Sandage, AR. 1959. M.N.R.A.S., 119, 255.
Freeman, K. 1988, private communication.

Haro, G. and Luyten, VVJ. ((1962), Bol.Obs.Tonantzintla y Tacubaya. 3,
No. 22, 37.

Harris, WE. and Racine, R. 1979. Ann.Rev.Astron.Astr0phys., 17, 241.
Hewett, RC. 1982. M.N.R.A.S., 201, 867.

Huchra, J.P., and Geller, M.J. 1982. Ap.J., 257, 423.

Lacey, C.G., and Ostriker, JP. 1985. Ap.J., 299, 633.

Phillip, A.G.D., and Sanduleak, N. (1968), 301. Obs. Tonantzintla. y Tacubaya,
4, No. 30, 253.

Pier, J.R. 1982. A.J., 87, 1515.
Pier, J.R. 1983. Ap.J.Suppl., 53, 791.
Pier, J.R. 1984. Ap.J., 281, 260.
Preston, G.W. 1988, private communication.
Ratnatunga, K.U. 1988, preprint.
Schechter, P. 1988, private communication.
Slettebak, A., and Brundage, R.K. 1971. A.J., 76, 338.

Sommer-Larsen, J., and Christensen, BR. 1986. M.N.R.A.S., 219, 537.

23
Sommer—Larsen, J ., and Christensen, RR. 1987, M. N.R.A.S. 225, 499
(SLC).
Sweigart, A.V. 1985, in Horizontal-Branch and U V—Brz'ght Stars, ed.
A.G. Davis Phillip (Schenectady: L. Davis Press), p. 3.
Wielen, R. 1988, in [AU Symposium 126; The Harlow Shapley Sympo-

sium on Globular Cluster Systems in Galaxies, ed. J. Gridlay

and A.G. Davis Phillip, (Dordrecht: Kluwer), p. 393.

STEVE P. DOINIDIS and TIMOTHY C. BEERS: Department of Physics and As-
tronomy, Michigan State University, East Lansing, MI 48824-1116.

24

TABLE 1

A PROPOSED GROUP OF FHB STARS ON PLATE CS 22936“

 

 

 

Velocity W K)
Star RA. (1950) Declination V B-V U—B DM (km/s) ( ) Comments
\m

279 18 58 09.4 -35 49 33 ......

281 18 56 33.4 —35 53 34 7 0.25 0.16 13.57 -—111 1.8

282 18 56 37.3 —36 03 14 7 0.02 -0.04 13.83 73 0.6 non—member
283 18 57 27.0 —36 06 18 0 0.11 0.14 13.62 -62 0.8

284 18 57 43.0 -36 06 32 6 0.02 0.03 13.62 —67 0.9

285 18 56 44.0 -36 10 51 8 0.11 0.15 13.70 201 0.6 non—member
286 18 57 44.6 -36 12 52 6 0.16 0.19 13.83 —98 0.7

287 18 57 30.2 -36 21 36 2 0.09 0.04 13.81 —60 0.7
290185911.9 —354754 ......

 

“Photometry from Pier (1982b). Radial velocities and K—line equivalent widths
from Pier (1983). Distance moduli assume the parameterization of Sommer—Larsen
and Christensen (1986).

25

FIGURE CAPTION 5

Figure 1 - (a) Mean Monte Carlo estimator of the two-point angular corre-
lation of the FHB I catalog. The error bars represent the one-sigma scatter about
the mean. The solid lines represent the one-sigma spread in the correlation of the
random catalogs against themselves.

Figure 1 - (b) Mean Monte Carlo estimator of the two—point linear correlation
of the FHB I catalog, using coarse apparent magnitude bins-toassign distances. The
error bars represent the one-sigma scatter about the mean. The linesrepresent the
one-sigma spread in the correlation of random catalogs against themselves.

0(9)

 

o—I
u-T
—'l
.1
.1
1

rlllllllllllllllllllllll'IIII

 

i-o—i
i—o—r

lllllllllllllllllllllllllllll

 

 

 

: / -
:- 1
- L l I L L l l l I l l l l l I l l -
O 20 40 60 BO

separation (arcminutes)

Figure 1 (a) - Doinidis and Beers (1989)

a.)

 

-61rlllrrllir[lilllri
.5—
\e.
O
.4

 

Illl'llll'llllllllllllll

I
y-s
lllllllll

 

llllllllllllllllllllllll

_lllll

lllllllll

 

 

.I
N
O

LIlLIILIIIli'iIlIII

20 40 6O 80 100
separation (parsecs)

Figure 1 (b) - Doinidis and Beers (1989)

CHAPTER 4: THE EXTENSION OF THE HK SURVEY TO THE

NORTHERN GALACTIC HEMISPHERE

The extension of the HK Survey to include the northern galactic hemisphere
began in 1986 utilizing the Burrell Schmidt telescope at KPNO, the sister telescope
to the Curtis Schmidt at CTIO, and an obvious choice for the continuation of the
survey. Plates are acquired in the same manner as in the southern hemisphere,
namely, a 4° prism is used together with the HK filter, and the exposures are
taken on Kodak IIao plates which have been hypersensitized. The exposure time
is nominally 90 minutes, but the images are widened, meaning that nine successive
offsets are made in right ascension, each occurring after ten minutes. The resultant
spectra have a dispersion of 180 A /m at Ca II H and K. The purpose of widening
the exposures is to gain clarity for the spectral features by enlarging the spectral

image.

The survey plates are scanned visually for candidates, and the criteria for the
different candidate classes is summarized in Table 1. Of the candidates which are
identified, roughly 35% are metal-poor candidates (type MP, although a number of
those classed as type B and C are expected to be metal-poor as well), 50% are

28

29

TABLE 1. HK Survey Stellar Classifications.

 

 

 

CODE Spectral Criteria
C Continuous spectrum; no lines present
D Degenerate spectrum; extremely broad He
B B-type spectrum; very blue continuum, weak He
AB AB-type spectrum; blue continuum, strong He, no CaII K visible
A A-type spectrum; strong He, weak Ca II K visible
MP Metal-poor F- or G-type spectrum; extremely weak CaII H and K
AF Intermediate between A- and F-type spectrum; moderate He/CaII H,
weak to moderate Call K
L Late-type spectrum; red continuum
E Emission noticed in spectrum
P Peculiarities noticed in spectrum

 

 

 

30
FHB candiates (types AB and A), and the remaining 15% are stars whose spectral
features were indicative of other interesting types, such as degenerates. Positions for
the candidate stars are measured using the XY measuring engine at Case Western
Reserve University. SAO positional standard stars are identified for each plate, and
20 to 25 are used for the plate solution. The typical accuracy which results in the
estimated positions is roughly 2 arcseconds, based on the residuals of the calculated

positions of the standard stars against their true positions.

As the intention of the survey is to identify candidate stars with dispatch, the
relatively poor resolution of the spectra make follow-up measurements to provide
confirmation of the candidate classes and to provide detailed information for those
which are successfully identified. Central to this work are those stars which are FHB,
A type, and so the resulting discussion will be devoted to the process of confirming
candidates of this type and the characterisation of the physical properties of these

stars based on the follow-up observations.

To date, spectroscopic observations have only been made for those stars in the
southern HK Survey. Spectra for the candidate stars are taken over the wavelength
range 3700 A to 4500 A, and a detailed discussion of the observation and reduction
procedures are given in BPS I and BPS II. Roughly equal-numbers of the spectra
were taken with the reticon spectrograph (z 0.7 A FWHM resolution) and the 2D-
Frutti (z 1.2 A FWHM resolution) at the DuPont 2.5m telescope at Las Campanas.
The candidate classes MP, C and B are assigned the highest priority for spectra, and

so most of the FHB and A type stars which have been measured spectroscopically are

31
misidentified metal-poor candidates. Radial velocities are obtained from measures
of line center for prominent absorption features, typically 4 to 6 for each star, with

an external error in the determination of roughly 10 km s71.

Equivalent widths for prominent features (H7, H6 and HeI) are obtained over
a fixed bandwidth of 12 A. The width of the Call K feature, KP, is measured over
bandwidths of 6, 12 and 18 A, and the equivalent width which is determined is
as follows: If the measured width is less than 2 A, then the width from the 6 A
bandpass is adopted, if the width is between 2 and 6 A, then the 12 A bandpass
measure is used, and if the width is greater than 6 A, then the 18 A measure is

adopted. The wavelength bands employed for calculating these widths are given in

Table 2.

A generalized index, HP, is calculated from H6 and H7 according to

HP = 0.5H6 + 0.588H7 — 0.276

This comes from a linear fit of H7 vs. H6, which is given in Figure 2 and follows

the relation

H7 = 0.468 + 0.850H6

For each star, the breadth of the H6 feature at 20% below the level of the local
continuum, D”, and its depth relative to the continuum, RC is also measured,

with errors of order 10% and 15% respectively. As many of the spectra that have

32

TABLE 2. Line Index Wavelength Bands (A).

 

Line

Line Band

Blue Sideband

Red Sideband

 

 

CaII K6
CalI K12
CaII K18
H 6
H 7
HeI A4026
HeI A4388

HeI A4472

 

3930.7-3936.7

3927.7-3939.7

3924.7-3942.7

40958-41078

43345-43465

40202-40322

4381.9-4393.9

4465.5-4477.5

 

3903.0-3923.0

3903.0-3923.0

3903.0- 3923.0

4000.0-4020.0

4247.0-4267.0

4000.0-4020.0

4350.0-4370.0

4435.0-4455.0

 

4000.0-4020.0

4000.0-4020.0

4000.0-4020.0

4144.0-4164.0

4357.0-4377.0

4144.0-4164.0

4405.0-4425.0

4490.0-4510.0

 

 

33
been obtained have relatively low signal-to-noise ratios (5 _<_ S/N S 10), an equiv-
alent width of S 0.5 A was considered to be the minimum which could be reliably
measured. Stars which are either FHB or A type can be separated from other hot

star types on a diagram of D0,; vs. RC, as shown in Figure 3.

FHB stars, typically, exhibit broad, deep Balmer lines (RC _>_ 0.5 and D0,; 5 30
A), although they may sometimes have shallower lines that give an appearance of
a high luminosity or sdB type star. In general, KP must be less than 3 A, by
noting that as members of the halo population they will have [Fe/ H] S —1. Figure
4 shows that, for this limit, the probability of spuriously including stars of solar
abundance is negligible. Type A stars are characterized by strong, narrow Balmer
lines (RC 2 0.5 and D03 5 20 A) and the presence of metallic lines. If KP is greater
than 4 A, then a star is automatically assigned to this type, as Figure 4 indicates
that this star almost certainly has a solar metal abundance. This figure also shows,
however, that for 3 _<_ KP _<_ 4, there is a great deal of ambiguity in assigning a
star as FHB or A type on the basis of its metallicity. The estimated uncertainty
in a metallicity derived in this manner is on order 0.3 dex, large enough to require

determinations of additional parameters in order to correctly classify these stars.

When colours are not available, it is possible to estimate them using the syn-
thetic colours and line profiles of Kurucz (1979) and Buser and Kurucz (1978).
Figure 5 shows the theoretical unreddened (B — V) colours for stars with FHB and
A type gravities as a function of HP. The size of KP is temperature dependent for

A type stars, and for FHB stars, it can be seen from Figure 4 that a KP greater

34

 

 

 

 

 

O. Q,
7 - O FHB/A o c ~.;—-,,._-;e -
.:;. .0? i
. .‘° Sign: «j O
0 0th er T e s “it“???i'i‘f’
yp C Waggfke
he" 0 '3:.,:‘,.;
3'. is?
:3 tees o C) .—
— 3.? sensug O
Ct'5‘riv'llse"” .
. o‘._ 1’- Qv.
‘ '1‘ g'IQ‘:I '.
.) ‘5; _&,t3;.-r,-‘
‘ .w“ x? {it I
rq‘ffii' 740D
9 Arvin-7‘. s @
‘ayn..‘r\n1‘;-} 5‘7.
— 31;‘<.¢i'il.:;\.v‘~’.o -
’s .y.‘“::""" ‘ a,“
9" ‘t‘Ifc'. ”Lg-‘3
”first“. «(r-5'.»
as!“ ugh-,9? :1} =41}
“ Jug-3; 2‘2?" ‘
. 93",: "r'.',‘;:_L‘J;';‘-.15;)
F “ISIS-'- 5 :-".-'.'-"1°""e
4 - O °;.;-:‘r‘-"..‘-. “70$" . d
o I -.j . ‘ '\
"11’ “'6. .*!
./ "0'1“...
,4

O
98
o o O” '0 00 O O

 

 

 

Figure 2. A plot of H 7 vs. H6 for hot blue field stars. Reproduced from Beers et
aL (1991).

35

 

GO'I'I'I'I‘I‘I‘I'ITI‘

55 "’ I Q ""

O ,1

50 .. ..
O O '

45 _ Oo .. IV -

 

 

 

 

4o — O -
o
,0
35 - -
O O .' Q
a 30 " co 00 "
S O O 83 00 (9 Q
Q 25 - ' , -
o a' c0
0'0 OO %
s C '
o 00,}? m 63% O O
O O
O #8”
10 - Eco -

 

 

 

 

5 , -
o -
5 III
_5 7 I L I l L L L L I l I l
0.0 0.1 0.2 0.3 0.4 0.5 0.6 0.7 0.8 0.9 1.0
R c

Figure 3. A plot of Dog vs. Rc for a variety of hot stars in the halo. Region (I)
is occupied by white dwarfs, st and sdB stars, region (II) by st, sdB and high
luminosity stars, region (III) by F HB and A type stars, and region (IV) by white
dwarfs. Figure reproduced from Beers et al. (1991).

....- .

tili .1

 

 

36

 

Call K Width (Angstroms)

 

 

 

 

0 I l L l l l l

0.00 0.05 0.10 0.15 0.20 0.25 0.30 0.35 0.40

(B-V)

Figure 4. Theoretical relation for KP vs. unreddened (B - V), showing loci of
[Fe/H]: -0.5 (upper line), [Fe/H]: -1 (middle line) and [Fe/H]: -2 (lower line).

37

 

0.6 I

0.5 -

(B-V)

0.2 -

0.1 -

 

 

 

 

0.0 L

HP

Figure 5. Theoretical relation for (B - V) vs. HP for [Fe/H]=-1. The solid line
represents the relation for log g = 3 and the dashed line, log a = .

38
than 3 A is indicative of a (B — V) colour redward of zero. Thus, only the region of
the (B — V) vs. HP relation redward of zero needs to be considered, avoiding the
problem of the turnover that the function has at this colour. This mean relation is
used to determine an approximate (B — V) colour for these stars, by averaging the

derived colour for each of the gravities, which can aid in deciding its type.

Surface gravity can also be used to discriminate between the two types, as FHB
stars have gravities that are lower than those of main sequence stars. This can be
determined if colours are available for the stars in question. So-called ”Kiel Dia-
grams” can be constructed using the synthetic colours and linewidths. In essence,
given values of (B — V)o, (U — B); and HP define loci on a graph of log g vs. Geff,
where 9,,” = 5040/ Te f f and T,” is in Kelvin. The intersection of these three loci
define a unique value of (log g,9,ff) for the star, with an accuracy of roughly $0.03
in 03” and $0.3 dex in log g. In general, log g a: 3 is typical for FHB stars and
log g a: 4 is typical for A stars. Even these methods sometimes fail to resolve the
ambiguity for these stars, and high resolution spectroscopy is required for a clear

determination of their type.

CHAPTER 5: PHOTOMETRY OF A SAMPLE OF STARS FROM

THE NORTHERN HK SURVEY

Published in
The Publicgtigs of the Astronomical Society of the Pacific,

1990 December.

39

40

ABSTRACT

Photoelectric photometry is presented for a sample of 139 halo stars drawn
from an extension of the HK objective-prism survey of Beers, Preston and Schect-
man to the northern galactic hemisphere. The candidates for which photometry is
reported here were selected to span a wide range of types, but are dominated by
stars classified as type AB, A, or metal-poor (MP) .

1. Introduction

The HK objective-prism/interferencefilter survey of Beers, Preston and Schect-
man has been underway for over a decade from the southern hemisphere; a total
of 193 usable plates have been obtained with the Curtis Schmidt telescope at the
Cerro Tololo Inter-American Observatory. Recently, this survey has been extended
to the northern hemisphere, where plates are obtained with the Burrell Schmidt
telescope at Kitt Peak National Observatory. To date this effort has yielded 92
usable plates, primarily in the direction of the north galactic pole. As described in
detail elsewhere (Beers, Preston, and Shectman 1985), each plate is visually scanned
to identify candidate halo stars (in the range of apparent magnitude 11 S B S 16)
which exhibit a weak or absent CaII K line. Roughly thirty five percent of the
candidates are classified as likely metal-deficient stars with effective temperatures
typical of halo main-sequence turnoff stars or cooler giants. On order fifty five
percent of the candidates exhibit an He Balmer line which is considerably broader
than that of a turnoff star, and are classified as type AB (no Ca II K line visible) or
A (weak CaII K line visible). The remaining 10 percent of the candidates selected
include likely subdwarf B stars, degenerate stars, and stars exhibiting a wide variety

of peculiarities in their spectra.

Follow-up photometry and spectroscopy of this sample demonstrates that the
HK survey technique is successful in identifying large numbers of the most metal-
deficient stars known in the Galaxy (Beers, Preston, and Shectman 1990, hereafter
referred to as BPS II). Photometry of the candidate AB- and A-type stars indicates

41

that the vast majority of these objects are members of the field blue horizontal-
branch (hereafter referred to as BHB) population (Pier 1982; Preston, Shectman,
and Beers 1990a, hereafter referred to as FHB II). The catalog of field horizontal-
branch candidate stars is already very large, and promises to expand rapidly as the
survey continues. Beers, Preston, and Shectman (1988) presents a list of some 4400
AB- and A-type candidate stars. Preston, Shectman, and Beers (1990b, hereafter
referred to as FHB III) employ U BV photometry for a subset of this catalog to
infer the existence of a small, but significant, gradient in the mean color of BHB
stars as a function of distance from the galactic center. Additional data for BHB
stars in the north galactic hemisphere would clearly be very useful.

Originally, we had only intended to obtain broadband photometry and derive
metal abundance estimates for those stars whose CaII K line strength available from
digital spectroscopy was consistent with a metal abundance [Fe/ H] S —2.0. How-
ever, Beers ct al. (1990) show that the correlation between CaII K line strength and
[Fe/H] is sufficiently tight to allow metal abundance estimation up to [Fe/ H] = —1.0.
As eighty percent of the low metallicity candidate stars picked out in the HK survey
meet this abundance criterion, we can profitably use photometric measurements for
virtually all of the MP candidates. A two-color (U — B)o versus (B — V)o dia-
gram provides a useful gravity discriminant between metal-poor stars at the main-
sequence turnoff and stars with similar temperature but lower surface gravity. The
division between hot metal-poor stars and BHB or asymptotic giant-branch stars
is dificult to draw from moderate signal-to-noise spectroscopic observations alone,

but is trivial to determine with UBV photometry (see BPS II).

In this paper we present broadband U BV observations for 139 stars identi-
fied from the northern extension of the HK survey. In Section 2 we describe the

observation and reduction procedures that we employed. Section 3 presents our

results.

42

2. Observations and Reductions

2.1 Sample Selection

The program stars were selected from the six northern hemisphere HK survey
fields listed in Table 1. For each plate we list approximate equatorial coordinates
of the plate center, galactic coordinates, and the number of stars observed in each
field. Positions for each star, accurate to two arcseconds, were obtained using the
XY machine at Case Western Reserve University. Enlargements of the ob jective-
prism plates served as finding charts.

2.2 Photometric Observations

All observations were obtained during a six night run in February 1990 with
the #2-0.9m telescope at Kitt Peak National Observatory. The data were acquired
using the AFP2 pulse-counting aperture photometer system, incorporating a 1P21
photomultiplier tube and standard Johnson UBV filters. A solid CuSO4 blocking
filter was included for observations with the U filter. Each observation consisted
of five integrations on the star and two integrations of the background sky in each
filter. The sky region was obtained by a 30 arcsecond offset in right ascension to an
area covered in the field of view of the guider camera. Care was taken to ensure that
faint stars were absent from this area. Integration times were adjusted for each star
such that a S/N ratio of 100 was achieved, based on Poisson statistics. Individual
integrations were no longer than 30 seconds. The longest total integrations were 400
seconds (U passband) for the faintest of the program stars. Offset guiding was not
deemed necessary, as repeated checks with the TV guider system indicated that the
star being measured stayed well within the 15 arcsecond aperture at all times. The
data acquisition software incorporated into the AFP2 system was used to provide
a first-order reduction of the colors obtained, along with their approximate errors;

this allowed for an assessment of the photometric quality of the observations as they

were made.

Standards were chosen from the celestial equatorial list of Landolt (1983). In

43

Table 2 we list apparent magnitudes and colors (according to Landolt) for 28 stan-
dard stars with 8.08 S V 5 11.80 and —0.24 S (B — V) S 1.36. Over the course of
each night, 20 to 30 measurements of stars from this list were made. Several mea-
surements were carried out each night at large airmasses (X _>_ 1.5) to give weight
to the extinction solutions. Program stars were observed exclusively in the vicinity

of the meridian (X 5 1.3). In all, 151 observations were made of 139 program stars.

2.3 Reductions

Photometric reductions were obtained by use of the method described by Har-
ris, Fitzgerald and Reed (1981). At our request, C. Reed kindly provided the
FORTRAN code which performs the reductions outlined in that paper. The trans-
formation coeflicients to the standard UBV system were solved for using the full
set of standard star measurements, with nightly determinations of the zero points
and first-order extinction terms. These proved consistent with expected values.
Second-order terms were determined for the entire run, and in general represented
a small correction to the derived colours. In Table 2 we list our determinations
of computed residuals from fits to the standard stars. Internal rms errors in the
reduction are 0.011 mag in V, 0.009 mag in B — V and 0.015 mag in U — B, based
on the residuals of the standards used for the solution. Figure 1 is a plot of the
differences in computed and standard values of B — V and U -— B colors for the
standard stars as a function of the the standard B -— V color.

External errors may be estimated from the small number of repeat observations
(obtained on different nights) available for ten of the program stars. Absolute
differences in measured magnitudes and colors for repeat observations are presented
in Figure 2. The solid lines are locally weighted regression lines (lowess, Cleveland
and Devlin 1988) fit to the absolute differences in apparent magnitude and color as a
function of the mean apparent magnitude. If the residuals in our measurements are
consistent with draws from a normal distribution, then the range between successive
measurements of individual stars is directly related to the standard deviation of the

parent distribution (Pearson and Stephens 1964). The dashed lines are lowess lines

44

for the implied standard deviations. The mean range in measured V, B — V, and
U --B colors for these ten stars are 0.036 mag, 0.021 mag and 0.027 mag, respectively.
Based on the calculations of Pearson and Stephens, these imply equivalent standard
deviations of av = 0.018 mag, 034/ = 0.011 mag, and 011-3 = 0.014 mag. It is
dificult to be certain of any trends in the data from batches this small. Nevertheless,
the range in repeat measurements in the V band does seem to increase for fainter
stars, an effect that may be unavoidable with limited integration times on an 0.9m
telescope. The ranges in colors appear roughly stable over the available span of

apparent magnitude.

Seven of the stars which are reported here have also been measured by others,
and a comparison of the derived colours is given in Table 3. These data were
acquired from the SIMBAD database operated by CD8, and are largely consistent
with the measurements of this paper. One exception is the V magnitude reported
by Mermilliod and Mermilliod (1990) for BSI6027-061, which differs also from that
reported by Dahn et al.(1982).

 

45

3. Results

Data for the program stars are summarized in Table 4. Column (1) lists the
star. Equatorial and galactic coordinates are listed in columns (2)—(5). The derived
magnitude and colors for each program star are listed in columns (6)—(8). An
approximate reddening for the direction of the plate centers, taken from Burstein
and Heiles (1982), is listed in column (9). The original classification of each program
star, based on the appearance of its objective-prism spectrum, is listed in column
(10). A colon next to the classification code indicates that some doubt exists in the
classification. A de-reddened two-color diagram for the program stars is given in
Figure 3. Four stars with (B — V),, > 1.0, all of which were classified in the survey as
peculiar, have been excluded from this diagram. Plot symbols are coded to represent
the original prism survey classifications. Their size is roughly representative of the
external 20 errors. Included in this plot are polynomial fits for luminosity classes
lab and III (Fitzgerald 1970), and class V (Johnson 1966 and Fitzgerald 1970). The
blackbody line is a fit to the data of Arp (1961) and Mathews and Sandage (1963).
Horn the work of Pier (1982) and FHB II, we expect the majority of the stars with
(B — V)a < 0.35 to be members of the field blue horizontal-branch population.
Those stars redward of (B — V)o = 0.35 are expected to be predominantly low
metallicity stars.

Much work remains to be done. A comparison of Figure 3 with Figure 8 of
FHB II establishes that a photometrically similar set of candidate stars is isolated
by visual scans (by Beers) of HK objective-prism plates obtained with the Bur-
rell Schmidt telescope to candidates selected (by Preston) from HK survey plates
obtained from the Curtis Schmidt. A dedicated program 'of U BV photometry of
additional candidates would be extremely useful for detailed analysis of the proper-
ties of the galactic field horizontal-branch component along the lines of FHB III, as
well as for estimation of metal abundance as described in Beers et al. (1990). Philip
(1987) outlines an ambitious program of spectrophotometric observations of A—type

field horizontal-branch stars. Several of the AB- and A-type stars listed in Table 3

46

are as bright as V = 11—12, and thus might prove useful for such an application.

We would like to acknowledge the assistance of Bret Goodrich of the KPNO
staff for helpful guidance at the telescope, and P. Pesch of Case Western Reserve
University for providing access to the Case measuring engine. S.P.D. would like
to commend the KPNO ”First Response Medical Unit” for doughty late night as-
sistance during the run. This work received partial support from the National
Science Foundation via grants AST 86-17265 and AST 90-01376 to Michigan State

University.

 

47

REFERENCES

Arp, H.C. I961. Ap.J., 133, 874.
Beers, T.C., Preston, G.W. and Schectman, SA. 1985. A.J., 90, 2089.

Beers, T.C., Preston, G.W., and Shectman, SA. 1988. Ap.J.Suppl., 67,
461.

Beers, T.C., Preston, G.W., Shectman, SA, and Kage, J.A. 1990, A.J.,

in press.
Beers, T.C., Preston, G.W. and Schectman, SA. 1990, in preparation.
Burstein, D. and Heiles, C. 1982. A.J., 87, 1165.
Dahn, C.C. et al.1982. 11.1., 87, 419.
Fitzgerald, M.P. 1970. Astr.Ap., 4, 234.

Harris, W.E., Fitzgerald, M.P., and Reed, BC. 1981. Pub.A.S.P., 93,
507.

Johnson, H. 1966. Ann.Rev.Astr.Ap., 4, 193.

Klemola, A.R 1962. A.J., 67, 740.

Landolt, A.U. 1983. A.J., 88, 439.

Ljunggren, B. 1965. Arkiv.Astron., 3, 535.

Mathews, T.A., and Sandage, A. 1963. Ap.J., 138, 30.
Mermilliod, J .C., and Mermilliod, M. 1990, in preparation..
Pearson, ES, and Stephens, M.A. 1964. Biometrilca, 51, 484.

Philip A.G.D. 1987, in New Generation Small Telescopes, eds. D.S.
Hayes, R.M. Gener, and DR. Genet (Mesa, AZ: Fairborn Press),
p. 165.

Pier, J. 1982. A.J., 87, 1515.

Preston, G.W., Shectman, S.A., and Beers, T.C. 1990a, Ap.J.Suppl.,
submitted (FHB II) .

48
Preston, G.W., Shectman, S.A., and Beers, T.C. 1990b, Ap.J.Suppl.,
submitted (FHB III) .
Sletteback, A., Bahner, K. and Stock, J. 1961. Ap.J., 134, 195.
Upgren, AR. and Kerridge, S.J. 1973. Pub.A.S.P., 85, 721.

49

FIGURE CAPTION S

Figure 1 - Mean differences, computer minus standard values, for observations of
standard stars for (a) B — V and (b) U — B are plotted versus standard values of

B—V.

Figure 2 — Absolute differences in observed magnitudes and colors for program stars
with repeat measurements, plotted against their mean V magnitude (a)—(c), and

against their mean (B-V) color for (d)—(f).

Figure 3 -— De-reddened two-color (U — B)ovs.(B -— V)o diagram for the program
stars listed in Table 3. Polynomial fits to stars of luminosity classes Iab, III, and V
are indicated, as is a blackbody curve. The diameters of the symbols are roughly

equivalent to the external 20' errors in the measurements.

50

TABLE 1
Summary of Plate Fields for Program Stars

 

 

PLATE RA (1950) DEC (1950) l (°) 6 (°) Stars Observed
BS 15621 10 17.9 +25 10 208.0 +562 57
BS 15622 12 52.2 +25 11 321.3 +87.7 1
BS 15623 13 58.3 +25 10 83.8 +648 1
BS 15625 11 46.1 +25 12 218.3 +75.7 8
BS 16026 12 26.1 +30 12 182.4 +84.3 55
BS 16027 13 12.2 +30 11 63.2 +84.2 17

 

 

Summary of Standard Star Data

51

TABLE 2

 

Standard Values

Computed - Standard

 

STAR N V B—V U-B AV A(B—V) A(U—B)
(1) (2) (3) (4) (5) (6) (7) (8)
BD+2 2711 1 10.367 —o.162 —0.708 +0001 —0.008 +0024
BD+5 2468 4 9.348 —0.116 —0.560 +0019 +0013 +0.011
HD 84971 5 8.636 —0.159 —0.770 +0007 —0.008 +0000
HD 100340 4 10.117 —0242 —o.975 +0001 —0.004 +0001
HD 118246 1 8.089 —0.141 —0.636 —0015 —0027 —0008
SA 99—296 3 8.454 +1. 187 +1. 265 -o.003 —0004 —0.002
SA 99—358 12 9.605 +0 776 +0. 509 —0.008 +0006 —0.007
SA 99—367 9 11.149 +1005 +0829 -0.001 —0001 +0004
SA 99—408 5 9.807 +0407 +0 043 -—0.011 +0001 +0003
SA 99—438 8 9.399 —0.155 -0.719 —0.006 —0005 —0.010
SA 99—447 9 9.415 —0071 —0.217 +0001 +0005 +0008
SA 100—95 5 8.915 +0814 +0391 +0006 +0004 —0019
SA 100—241 5 10.140 +0156 +0102 +0001 —o000 +0006
SA 100—280 5 11.802 +0496 +0009 —0.002 —0.012 +0012
SA 100—606 7 8.641 +0052 +0125 -—0.004 +0.006 -0.036
SA 101-281 5 11.579 +0814 +0435 +0006 —0.003 +0014
SA 101—363 6 9.871 +0262 +0121 —0003 +0007 +0007
SA 102—58 6 9.380 +0060 +0021 —o.004 +0008 —0.008
SA 102—620 5 10.067 +1087 +1013 -0.006 +0010 +0008
SA 102-625 6 8.890 +0552 +0035 +0009 —0007 —0.004
SA 103—302 4 9.862 +0. 369 -0.058 +0002 —0005 +0009
SA 103—462 5 10.111 +0 564 +0089 +0015 —0.008 +0.011
SA 103—526 8 10.903 +1089 +0941 +0005 -0.002 —0.003
SA 104—337 6 11.207 +0768 +0336 +0001 -0000 +0015
SA 105—448 5 9.176 +0249 +0037 +0007 —0.006 —0.004
SA 106—485 1 9.484 +0 380 —0. 039 —o.015 +0004 —0004
SA 106—834 1 9.088 +0701 +0292 —0.012 —0003 +0018
SA 106—1024 1 11.594 +0332 +0 085 +0015 +0020 +0009

 

 

52

TABLE 3

Comparison with Previous Photometric Measurements of Program Stars

 

 

This Paper Other Work
STAR V B — V U — B V (B — V) (U — B) Source

(1) (2) (3) (4) (5) (6) (7) (8)
15625-002 10.68 +1.51 +1.19 10.68 +1.52 +1.23 6
027 11.54 +0.30 +0.06 11.60 +0.26 +0.06 2
16026-010 10.10 +0.33 —0.03 10.08 +0.36 . . . 3
017 10.65 +1.44 +1.24 10.63 +1.46 +1.25 4
10.62 +1.42 +1.27 6
056 11.00 +0.17 +0.11 11.04 +0.19 . . . 3
11.05 +0.16 +0.09 2
11.00 +0.17 +0.07 1
16027-061 12.66 —0.14 —1.20 12.68 -0.12 —1.17 5
12.86 —0.10 —1.14 6
070 13.75 +0.04 +0.02 13.79 +0.07 +0.02 6

 

 

999939!"

Slettebak et.al.(1961).
Klemola (1962).

Ljunggren (1965).
Upgren and Kerridge (1973).
Dahn et.al.(1982).
Mermilliod and Mermilliod (1990).

53

TABLE 4

Summary of Photometric Data for Program Stars

 

 

STAR RA (1950) DEC (1950) l(°) b (°) V B - V U — B EB-V CLASS
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10)
15621— 1 10 27 29.8 +26 16 09 206.8 +58.5 12.38 0.66 0.22 0.00 P
2 10 26 24.4 +26 14 18 206.8 +58.3 11.82 0.16 0.11 AB
7 10 25 10.3 +28 03 26 203.3 +58.3 12.13 0.70 0.23 MP:
8 10 24 50.4 +28 03 26 203.3 +58.2 11.24 -0.03 -0.05 AB
9 10 25 13.4 +27 47 58 203.8 +58.3 13.59 0.25 0.07 A/MP
10 10 25 19.1 +26 48 11 205.6 +58.1 13.75 0.65 0.21 MP/A
11 10 23 52.8 +25 59 38 207.0 +57.7 13.98 0.65 0.11 MP
12 10 24 30.4 +25 37 33 207.8 +57.7 12.27 0.22 0.09 MP/A
14 10 23 43.7 +24 06 04 210.5 +57.2 13.70 0.68 0.16 P
15 10 23 33.6 +23 45 53 211.0 +57.1 12.26 0.14 0.16 AB/A
16 10 23 30.5 +28 02 43 203.2 +57.9 13.36 0.65 0.02 MP
17 10 22 10.0 +27 55 06 203.4 +57.6 11.09 0.09 0.10 A
18 10 21 31.8 +26 37 46 205.7 +57.3 14.19 0.67 0.18 MP:
19 10 23 07.7 +26 16 11 206.5 +57.6 11.93 0.16 0.11 A
20 10 21 32.1 +24 23 05 209.7 +56.8 10.96 -0.05 -0.06 AB
21 10 22 22.0 +23 56 09 210.6 +56.9 12.49 0.19 0.12 A/MP
22 10 19 07.1 +24 06 06 210.0 +56.2 12.85 0.30 0.00 MP/A
23 10 19 08.0 +24 31 20 209.3 +56.3 13.30 0.48 -0.05 MP
24 10 19 21.2 +25 42 08 207.2 +56.6 14.78 0.38 -0.27 C
26 10 17 14.7 +24 11 51 209.6 +55.8 13.65 0.62 0.02 MP
27 10 18 19.7 +23 14 33 211.4 +55.8 12.70 -O.26 -1.05 AB
28 10 15 38.9 +24 37 07 208.8 +55.6 13.89 0.61 -0.02 A:
29 10 15 11.9 +25 15 47 207.6 +55.6 13.94 0.65 0.01 MP:
30 10 15 50.7 +26 00 48 206.4 +55.9 13.41 0.69 0.22 MP:
31 10 16 50.7 +26 38 11 205.3 +56.2 14.60 0.31 0.10 MP/A
32 10 16 27.7 +26 42 31 205.2 +56.2 14.83 0.04 0.17 AB
33 10 13 20.4 +26 10 29 205.9 +554 13.73 0.61 -0.08 MP
34 10 13 21.5 +26 04 27 206.1 +55.4 13.65 0.77 0.43 P
36 10 12 51.9 +25 20 54 207.3 +55.1 11.67 1.31 1.26 P
37 10 13 17.5 +24 59 15 207.9 +55.1 11.64 0.23 0.02 A
38 10 11 30.9 +23 49 12 209.7 +54.5 13.39 0.49 -0.09 MP:
39 10 11 02.6 +25 33 18 206.8 +54.8 13.60 0.06 0.13 AB
40 10 09 42.1 +27 32 30 203.3 +54.8 13.81 0.20 -0.03 A
41 10 08 39.5 +23 13 56 210.4 +53.7 12.60 0.22 0.05 A
42 10 07 03.7 +24 12 59 208.7 +53.6 13.70 0.56 0.01 MP
43 10 06 38.9 +25 04 50 207.2 +53.7 14.41 0.05 0.07 A:

 

54

TABLE 4 (Continued)

 

 

STAR RA (1950) DEC (1950) l(°) 0 (°) V B — V U — B E3_V CLASS
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10)
15621— 45 10 06 27.5 +26 16 26 205.2 +53.9 14.00 0.62 0.07 MP
47 10 06 46.6 +26 14 17 205.3 +53.9 13.63 0.87 0.37 C:
48 10 14 18.0 +24 30 50 208.8 +552 12.20 0.26 0.08 MP/A
49 10 06 39.0 +25 37 02 206.3 +53.8 13.22 0.25 0.13 AB
50 10 10 04.4 +23 38 02 209.9 +54.1 13.92 0.58 0.01 MP:
51 10 14 32.8 +23 18 16 210.9 +55.0 13.97 0.45 0.09 A
52 10 14 47.1 +23 13 33 211.1 +55.0 13.42 0.44 -0.10 MP
53 10 17 15.9 +23 09 05 211.4 +55.6 13.59 0.25 0.04 A
54 10 21 04.3 +24 10 40 210.1 +56.7 14.13 0.46 -0.26 A/MP
55 10 23 04.5 +24 11 01 210.2 +57.1 13.30 0.41 -0.14 MP
58 10 26 50.6 +23 40 02 211.6 +57.8 14.38 0.35 -0.07 MP
63 10 27 26.7 +26 37 21 206.1 +58.6 11.08 0.99 0.82 P
68 10 21 25.3 +23 05 13 212.0 +56.5 10.32 0.39 -0.01 MP
69 10 19 06.4 +26 22 05 206.0 +56.7 13.61 0.63 0.11 MP
70 10 19 07.9 +27 26 28 204.1 +56.9 12.61 0.37 -0.07 MP
71 10 12 44.2 +27 15 19 204.0 +55.4 13.02 0.42 -0.03 MP
72 10 14 29.2 +26 05 13 206.1 +55.6 12.80 0.30 0.10 MPz/A:
73 10 10 30.5 +24 19 59 208.8 +54.4 13.53 0.42 -0.21 0.00 MP
74 10 10 05.6 +27 27 30 203.5 +54.9 13.69 0.62 0.21 MP:
76 10 08 02.9 +26 28 26 205.0 +54.3 12.92 0.44 -0.02 MP
77 10 10 11.7 +25 04 42 207.5 +54.5 13.87 0.43 -0.17 MP
15622— 36 13 03 49.7 +28 20 21 49.9 +86.6 11.97 0.96 0.80 0.01 A:
15623— 1 14 06 11.8 +23 35 02 24.9 +722 13.75 0.14 0.19 0.00 AB
15625- 2 11 39 33.5 +26 59 11 210.5 +74.6 10.68 1.51 1.19 0.01 MP:
15 11 54 27.0 +25 43 03 217.8 +77.7 13.30 0.52 -0.08 MP
17 11 53 50.0 +24 44 53 222.0 +77.3 12.86 0.67 0.00 MP:
18 11 57 11.4 +25 22 46 219.9 +782 13.77 0.50 -0.15 MP
23 11 47 59.7 +28 01 07 206.8 +76.5 13.13 0.49 -0.04 MP
24 11 50 00.8 +23 36 41 225.6 +762 11.81 0.48 -0.06 MP
26 11 49 14.9 +26 38 37 212.8 +76.7 11.07 0.30 0.08 MP
27 11 53 07.6 +23 16 39 228.0 +76.7 11.54 0.30 0.06 MP
16026- 2 12 14 52.1 +32 06 50 178.5 +812 13.67 0.67 0.06 0.01 MP:

 

55

TABLE 4 (Continued)

 

 

STAR RA (1950) DEC (1950) l(°) 0 (°) V B — V U — B 153.1; CLASS
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10)
16026— 3 12 14 40.4 +31 45 51 180.6 +81.4 13.94 0.63 0.04 MP
4 12 15 12.8 +31 34 36 181.4 +81.5 13.66 0.66 0.06 MP:
5 12 14 34.6 +31 02 25 185.1 +81.7 14.15 0.70 0.18 MP
6 12 14 26.5 +30 59 24 185.5 +81.7 14.35 0.35 -0.19 MP
7 12 14 21.5 +30 37 54 187.8 +81.8 13.89 0.68 0.14 MP
8 12 15 35.0 +30 34 04 187.6 +82.0 14.80 0.08 0.18 AB
9 12 15 16.8 +30 31 47 188.0 +82.0 14.15 0.60 0.03 MP:
10 12 14 15.4 +29 51 49 193.0 +82.0 10.10 0.33 -0.03 MP
11 12 14 18.8 +29 12 43 197.6 +822 13.56 0.07 0.13 AB
12 12 15 04.0 +28 42 37 201.1 +82.4 13.94 0.58 0.02 MP:
14 12 13 55.6 +28 07 19 205.6 +822 14.72 0.21 0.13 MP/A
15 12 17 26.9 +27 54 06 207.1 +83.0 13.41 0.62 0.07 MP
16 12 15 50.8 +28 17 56 204.1 +82.6 14.17 0.61 -0.05 AB:
17 12 16 54.0 +28 39 31 201.1 +82.8 10.65 1.44 1.24 P
18 12 17 38.6 +29 01 06 198.0 +82.9 13.86 0.47 -0.21 MP
19 12 15 46.9 +30 09 39 190.3 +822 14.20 0.52 -0.19 MP
20 12 16 36.1 +32 07 20 177.2 +81.5 12.34 1.49 1.19 MP:
21 12 20 03.0 +32 16 05 173.5 +82.1 13.91 0.75 0.40 MP
23 12 19 27.4 +31 23 19 179.4 +82.4 12.80 0.17 0.05 A
24 12 19 51.5 +31 04 12 181.3 +82.7 14.84 0.42 -0.01 C
26 12 19 32.9 +28 20 36 203.0 +83.4 13.95 0.08 -0.04 AB
27 12 19 21.4 +28 02 11 205.7 +83.4 13.66 0.76 0.36 MP
28 12 20 32.4 +27 43 53 208.3 +83.7 13.68 0.01 0.08 AB
29 12 21 22.2 +31 43 53 175.5 +82.6 14.13 0.55 -0.03 MP:
30 12 23 36.2 +31 38 31 174.0 +83.0 14.33 0.51 -0.06 MP
32 12 23 35.4 +31 16 48 176.5 +832 13.57 0.77 0.38 MP:
33 12 24 33.5 +29 38 17 188.9 +84.2 14.46 0.55 -0.02 MP
35 12 24 55.9 +28 01 30 204.9 +84.6 13.76 0.38 -0.12 A/MP
36 12 26 14.2 +28 37 11 198.0 +84.8 13.91 -0.12 -028 AB
38 12 27 04.2 +29 50 17 184.8 +84.6 12.94 0.36 -0.05 MP
40 12 27 56.6 +31 11 22 172.2 +84.0 13.35 0.40 -0.19 MP
41 12 28 42.4 +31 08 01 171.6 +84.2 14.05 0.11 0.13 AB
42 12 28 49.8 +30 50 25 173.8 +84.4 13.82 0.58 0.02 MP
43 12 28 32.5 +30 41 15 175.5 +84.5 14.22 0.47 0.02 MP
47 12 29 13.3 +28 28 46 198.0 +85.5 14.08 0.67 0.16 MP:
48 12 30 17.0 +28 27 33 197.5 +85.7 13.33 0.90 0.54 0.01 MP:

 

56

 

 

TABLE 4 (Continued)
STAR RA (1950) DEC (1950) l (°) b (°) V B — V U — B E3_v CLASS
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10)
16026— 49 12 30 06.7 +28 32 28 196.6 +85.7 12.30 1.27 1.10 MP:
50 12 31 25.2 +28 33 03 195.5 +86.0 13.62 0.59 0.00 MP
54 12 30 35.9 +29 54 00 180.1 +85.3 14.34 0.54 -0.09 MP
55 12 29 53.6 +30 06 46 178.9 +85.0 14.48 0.59 -0.01 MP
56 12 31 45.1 +30 56 21 168.7 +84.8 11.00 0.17 0.11 A
57 12 30 24.5 +31 18 35 167.9 +84.4 14.40 0.43 -0.14 A/ MP
58 12 29 55.8 +32 09 52 162.8 +83.7 13.46 0.66 0.12 A
59 12 34 05.7 +32 20 02 155.3 +84.1 14.48 0.22 -0.02 A13
61 12 32 54.2 +31 34 39 162.0 +84.5 14.25 0.15 0.07 AB
62 12 32 06.1 +31 01 09 167.5 +84.8 13.53 0.60 0.04 MP:
63 12 32 02.6 +30 52 52 168.8 +84.9 12.81 1.17 1.00 MP:
64 12 32 17.3 +30 40 16 170.2 +85.1 13.80 0.54 0.01 MP
65 12 32 30.3 +30 32 07 171.0 +852 13.42 0.66 0.11 MP:
66 12 33 37.9 +27 56 49 203.0 +86.6 14.12 0.37 -0.10 MP
67 12 33 53.4 +27 33 04 209.5 +86.6 13.39 -0.01 0.04 AB
72 12 35 10.4 +30 49 35 163.6 +85.4 13.54 0.62 0.08 MP
73 12 36 31.8 +30 59 14 159.4 +85.5 13.38 0.33 -0.08 A/MP
74 12 34 52.0 +32 01 50 155.7 +84.4 14.45 0.56 -0.11 MP
16027— 1 13 03 52.2 +32 58 21 93.9 +83.6 12.82 0.33 0.01 0.00 A/MP
3 13 04 53.5 +32 33 26 90.3 +83.8 13.77 0.43 -0.16 MP
9 13 03 05.2 +31 13 24 85.1 +85.1 14.11 0.65 0.15 E
15 13 03 38.8 +30 31 39 78.2 +85.5 14.06 0.53 -0.05 MP
28 13 06 44.3 +29 15 48 59.6 +85.7 12.80 0.60 0.12 A/MP
35 13 06 21.0 +33 08 08 90.9 +832 11.69 -0.01 0.06 AB
38 13 08 28.2 +32 44 11 85.9 +832 11.73 -0.01 0.06 AB
39 13 08 56.0 +32 16 53 82.7” +83.5 13.79 0.64 0.12 MP:
49 13 10 05.2 +30 37 10 69.2 +84.4 13.26 0.39 0.04 A
53 13 10 52.5 +31 37 51 76.0 +83.6 14.60 -0.13 -0.55 AB
56 13 10 38.1 +32 53 56 83.9 +82.8 13.92 0.46 -0.19 MP
57 13 13 53.3 +32 12 41 76.2 +82.8 14.59 0.00 0.05 A
59 13 12 22.1 +31 47 00 75.2 +83.3 13.94 0.06 0.03 AB
61 13 14 00.3 +29 21 44 54.1 +84.2 12.66 -0.14 -1.20 C
63 13 12 21.4 +28 38 59 48.0 +84.7 15.00 0.36 -0.20 MP
70 13 17 21.9 +29 39 28 54.6 +83.4 13.75 0.04 0.02 AB
84 13 22 42.7 +30 01 57 54.6 +822 14.12 -0.04 0.11 AB

 

 

57

CLASS CODES: AB— type AB; A- type A; MP— metal poor candidate; C- contin-

uous; P- peculiar; E— emission

DBLTAJU-B)

DBL’I‘AJB-V)

 

 

 

 

 

 

 

 

Figure 1 (a) and (b) - Doinidis and Beers 1990

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CHAPTER 6: A DETERMINATION OF THE ROTATIONAL

VELOCITY OF A SYSTEM OF FIELD STARS

The application of halo samples to kinematic models for the galaxy requires
a knowledge of seven things for each object in the sample: the three components
of both its position and velocity, and the form of the potential through which it
moves. Of these, only four can be directly observed, namely, position and heliocen-
tric radial velocity. If the number of objects in the sample is small, then simplifying
assumptions can be made with regard to the symmetry of the potential or the distri-
butions of position and velocity, in order to constrain the rest. With larger samples,
these constraints can be made statistically. It is also preferable, in any case, to use
samples which are uniform in general properties (e.g., stellar type and metallicity).
This ensures that the objects comprising the sample share other physical proper-
ties and the same evolutionary history. It is a relatively straightforward exercise to
determine the rotational velocity of a system of stars, and this determination can
be applied to important problems in galactic structure and evolution. The interface
between the disk Population I and the halo Population II must exist at some point,
for example, but it is not completely clear as to where one may draw a distinction

between the two, or even whether they are truly discrete. Eggen, Lyndon-Bell and

62

63
Sandage (1962), based on a correlation of rotational velocity with abundance, argue
that the galaxy underwent a rapid collapse to its present form, with progressive

enrichment of the material in the disk from young halo stars.

Also using metal abundance and kinematics as a guide, Gilmore et al. (1989)
review the evidence for a third component of the galactic system, the so-called
”thick” or ”extended” disk, which has high rotational velocity (but less than that
of the disk) and intermediate metal abundance, peaking at [Fe/H] z -—0.6 and a
characteristic scale height 2 a: 1.4 kpc, and which presumably formed during one or
more intermediate dissipative stages of the galaxy’s collapse. Norris (1986), how-
ever, using a large non-kinematic sample of stars in the solar neighbourhood, finds a
sharp transition between halo and disk kinematics at roughly —1.2 S [Fe/ H] S —1.6,
suggesting therefore that the disk system is decoupled from the halo, and moreover,
that for [Fe/ H] 2 —1.2, there is a marked increase in rotation that is smooth over the
region occupied by the ”extended disk.” It should also be noted that Norris differs
in his definition of the ”extended” disk, maintaining that it should be considered a

part of the disk rather than a separate component of the galaxy.

A sample of 592 stars, drawn from the southern hemisphere HK Survey, is
used to calculate the rotational velocity of the halo. Horn this sample, some 507
have metallicities which can be determined from the relations in Figure 4, and the
analysis which follows will include only these stars. There are two stellar types in
the sample: 269 FHB stars, and 238 stars with spectra similar to those of normal A

type main sequence stars. The stars have been separated according to the criteria

64
outlined in Chapter 4, with the additional criterion that stars with [Fe/ H] _>_ —0.75
are considered A stars, and those with [Fe/H] S -0.75 are FHB, in order to clearly
separate these two sub-samples. While this division is arbitrary, it serves to divide
the sample into halo and disk components, where the disk component (the A type

stars) is so named because of their assumed young age.

Lance (1988) found no systematic motion for a polar sample of these apparently
normal A type stars, and suggests that their young ages require that they had
been formed in situ, rather than had been ejected from the disk. She therefore
postulates that they were formed due to a collision or merger between the galaxy
and a satellite. It should be possible to test this hypothesis with this sample, as
the A type stars here should show no evidence of systemic rotation, or, at the very
least, rotation which is markedly different than that of the disk, ”extended disk” or

halo populations.

For a Cartesian coordinate frame centered on the sun, where the galactic center
is along the negative X-direction, the galactic rotation in the positive Y-direction,
and the NGP in the positive Z-direction, the distributions of this sample in the
Z-X and Y-X planes are given in Figure 6. The distributions of these stars in the
Z-direction over the range in their metallicities are given in Figure 7. It can be seen

that, over this metallicity range, both components of the. sample are distributed

uniformly.

The pertinent data for the sample are presented in Table 3. The star identifi-

cations (Column (1)) are composed of the plate number from the Curtis Schmidt

65

 

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Figure 7. Distance above the galactic plane, Z, vs. abundance for the F HB stars
(open circles) and A type stars (filled circles). The dashed lines represent a typical
distance for each of these types near the magnitude limit of the HK Survey.

67
plate log and an individual number for that star. Columns (2) through (5) give the
1950.0 coordinates for each star, and Columns (6) through (8) give the measured
UBV colours (where available). Column (9) gives the reddening for the direction of
each plate center, which was taken from Burstein and Heiles (1982), or from Preston
(1990) when available. In Column (10), (B — V) estimates from synthetic colours
are given for those stars without photometry. Column (11) gives the measured
heliocentric radial velocity for the star, Column (12) gives the detector used for the
spectrum (F = 2D-I'H'utti, R = Reticon), and Columns (13) through (16) give the
spectral parameters discussed in Chapter 4. Column (16) lists the type given to
that star from its spectral appearance of the plate (see Table 1), and Column (17)
give the final type which was assigned to the star. Columns (19) and (20) give the
absolute magnitude and heliocentric distance (in parsecs) for each star, and Column
(21) the derived metal abundance, again following the method outlined in Chapter
4. Stars which are marked with an asterisk are those for which KP was below the
measurable threshold, and so were not included in the abundance determination.
A number of stars are given a metallicity of ”< —2.00,” and these stars are those
whose measured values of KP put them considerably below the [Fe/ H] = —2 line of
Figure 4. It is assumed that they are of lower metal abundance, but this cannot be

confirmed without higher resolution spectra.

68

 

 

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88.8 8888 88.8 < 888 88.8 8.8 88.8. 8.8 8 .8 88.8 88.8 8.88. 8.8 888888- .88: .8 88

88.8- 8888 88.8 < 82 88.8 88.8 88... 88.8 8 88.- .88 88.8 8.8.1 8.8 888.88- 8.8.88.8 88

88..- 8888 8.8 8..... o 88.8 88.8 ...8. 88.8 8 88- 88.8 88.8 8.88. 8.8 .. 8888- 8..888.8 .8

88.8 8888 88.8 < .888 88.8 88.8 88.8 88.8 8 8- 88.8 88.8 8.88- 8.8 88 .8 88- 8.8 8. .8 8

88.8- 8888 8.8 88.8 o 88.8 88.8 88.8. 88.8 8 8. 8.8 88.8 8.88. 8.. 888888- 88. 88 .8 88

88.8 888 88.8 8. 888 88.8 88.8 88.8. 88.8 8 .8- 88.8 88.8 .88. 8.. 888888. 8.8 88 .8 88

88..- 888 88.8 88.8 < 88.. 88.8 88.8 88.8 8 88 8.8 88.8 8.8.1 8.8 888888- 88. 88 .8 88

88.8 8.88 88.8 8. o 88.8 88.8 8.8. 88.8 8 8- 88.8 88.8 8.8. 8.8 8888.81 8.88 8888 88

88.8 88.8 88.8 < 888 88.8 88.8 88.8. 88.8 8 8.1 88.8 88.8 8.8. 8.8 888888- 8.88 8888 88

88..- 8888 88.8 888 8 88.8 .88 88.88 8.8 8 8- 8.8 88.8 8..... 8.. 888888. 88888.8 8.

.8..- 8888 88.8 88.8 8 .88 88.8 88.8 .88 8 8- 8.8 88.8 8..? 8.. 8.88.8- 8.88.8 .. 88888

.88- .888 8.8 88.8 888. 88.. 88.8 88.8 88.8 8 88.- 888 88.8 8.88- 8.888 88 88 88- 8888888 88

88..- 888 8.8 888 888 8.. 88.8 88.8 88.8 8 88 88.8 8.8 ..8- 8.888 88888. 8.88 8888 88 .8888

8.8 .888 .8: 88.8 88.8 88.8 88.8 .8: .8: .8: ..: .8.. .88 .88 .88 .88 .88 .88 .88 .88 3
88. 8.8.8 8.. 888:. 888.5 88 88 28 .88 .88 88> 88 8-88 8:8 8:8 8 8 8 88.88.88 8888

 

“‘

88288888 .8 8.88..

87

 

 

 

88.88 88.8 88.8 8. H888 88.8 88.8 8.8. 88.8 8 88 88.8 88.8 8.88: 8.8.8 888 .8- 8888. 88 88
88.88 88.8 88.8 < .888 88.8 88.8 88.8. .88 8 88 88.8 88.8 ..88. 8.8.8 88 .88.- 88.8888 88
88.8. 8.88 88.8 < 888 88.8 88.8 ...8. 88.8 8 8- 88.8 88.8 8.8.: 8.8.8 88888.- 8888888 88
88..- 8888 8.8 8.88 8.8 8.. 88.8 88.8. 88.8 8 8 8.8 88.8 8.88. 8.8.8 8.888.- 88:88 88
88.8 8..8 88.8 8. 8.8 88.8 88.8 8.8. 88.8 8 ..- 88.8 88.8 8.88. 8.8.8 8.8 8.- 88.8888 88
8.8-v 8.88 88.8 888 8.8 88.8 88.88 88.8 8 88. 88.8 88.8 8.88. 8.8.8 8888.- 8888888 88
88.88 8888 88.8 < 8: 88.8 88.8 88.8. 8.8 8 88 88.8 88.8 8.8.: 8.8.8 8888.. 88.8888 88
88.88 8.88 88.8 8. .8: 8.8 88.8 88... 8.8 8 8. 88.8 88.8 8.88. 8.8.8 88888.- 8888888 88
88.8- 8888 88.8 8. 82>. 88.8 88.8 88.8. 88.8 8 88- 88.8 88.8 8.88. 8.8.8 .8888- 8888888 .8
88.8 8888 8.8 888 ”82 8.8 88.8 88.8. 88.8 8 .8- .88 88.8 8.8.: 8.8.8 .8888- 8888888 8 -8888
..8 .888 88.8 .8: .8: .8: .8.. .8: .8: .8: ..: 88.8 .88 .88 E .88 .88 .88 .88 .88 E
88. 88.8 888 888... 888.5 88 Q8 28 8.. ..88 88> 88 8-8.8 8:8 8:8 8 8 8 888888.888 8.88

 

88888888 .8 8.88..

88
Absolute magnitudes for these stars are determined according to the relations
given in Figure 7 of Preston (1990). For the FHB stars, absolute magnitudes are

assigned according to

(B — V) g 0.02 : Mv = 2.22
0.02 < (B - V) < 0.20 : My = 1.04 — 4.423(3 — V)2 + 17.74(B — V)3 (2)

(B — V) _>_ 0.20 : My = 0.60

and the A stars according to

Mv = 0.602 + 11.07(B — V) — 15.843(B — V)'-’. (3)

The method followed to determine the rotational velocity of the halo is that
given in F W80, and will be outlined here. In a galactocentric frame, let 129 be the
velocity of the local standard of rest (in the direction of the galactic rotation), 1&pr
the expansion velocity of the ith star, and v,“ the systemic rotational velocity. Us-
ing the quantities defined in Figure 8, then the observed heliocentric radial velocity

of the ith star, v°,.- is given by

vo,.- = vm cos w- + v8”- cos 05; + 1288,; — vo cos »\.~ (4)

so that the first two terms arise due to the systemic rotation and expansion of the

stellar system, the third from the peculiar heliocentric velocity of the ith star

89

ROTATION
AXIS

1
D‘\\

\

\ \ GLOBULAR
CLUSTER

 

   

 

GALACTIC \ \
CENTRE

Figure 8. Definitions of the geometric and physical parameters used in the FW80
method (Figure reproduced from FW80).

90
with respect to this system, and the last from the sun’s motion around the galac-
tic center. The projection factors cos z/u, cos ()5; and cos A,- can be related to the

measured heliocentric latitude and longditude by

cos A; = cos b.- sin I.- (5)

cos t/J,‘ = R6) cos b; sin MR? cos2 1),-sin2 I, + (R3 — R,- cos b,- cos 1,)2]'1/2 (6)

cos gt, = R,- - R9 cos b,- sin (JR? + R2) — 2R5RG cos 1, cos b,-]-1/2 (7)

where R3 is the distance from the sun to the galactic center (assumed here to be 8

kpc) and R,- is the derived heliocentric distance to the ith star.

In the simplest case, the expansion velocity can be considered either a constant,
or decreasing at a uniform rate 6 = van/r5, where r.- is the galactocentric distance
to the ith star. For this calculation, a constant value for v3,”- is assumed, and
moreover, it is assumed that it is small enough to be neglected. Even with this
assumption, one is faced with the problem of constraining both um, and up“; in
equation (2). To first order, this can be accomplished by solving for um while

neglecting vpec,.-, and then using this result to determine an rms estimate of up“,

If equation (2) is multiplied by cos 1b,- and summed over all stars, then one gets,

to first approximation,

91

v _ 2 cos tin-(veg + ”0 cos Ai)
mt" _ 2 cos2 $5

 

(8)

where 0,0,, is the rotational velocity assuming no expansion or peculiar velocities,

and, finally,

 

 

 

v _ 23 cos ¢,(v°,,- + 129 cos A.) (no, (9)
rat _ 2 cos'2 #2.- [2 cos2 ¢,]1/2
where (no, is the rms observed value of vpec,,- obtained from
0,0, = 2(v0fi. + ”6) C08 A! — ”708,. COS ¢!)2 1/2 (10)

N—l

where N is the number of stars in the sample, and with an assumed error for (no,

of (0'10,/2N)1/2.

The sample is divided into metallicity bins for the calculation. The results of
the calculations for two such sets of bins are given in Tables 4 and 5. Figure 9 gives
the results for the first set of bins (from Table 4), and the transition from halo to
disk rotation can be seen in roughly the same area ([Fe/ H] a: —1.2) found by Norris
(1986) and Norris and Ryan (1989). For larger bin sizes, this transition is not as
sharp, as can be seen in Figure 10 for the bins of Table .5, although Norris and
Ryan maintain that their result does not differ when the number of bins that they
use is reduced by a factor of up to two. Clearly, the results presented here have a
noticable dependence on bin number, and this dependence comes from the range in

[Fe/H] covered by the bin more than the number of objects per bin. It would

92

TABLE 4. FW80 Analysis Results.

 

 

 

([Fe/Hl) cure/a] N vm 0.- mo. (disp)zo. I ”rot/aloe I
(km 5") (km 8-1) (km 8”) (km S“)

(1) (2) (3) (4) (5) (6) (7) (8)
-1.98 0.17 40 -29 34 106 12 0.27
-1.65 0.08 40 49 24 80 9 0.61
-1.40 0.08 40 53 29 103 11 0.51
-1.17 0.06 42 38 30 98 11 0.38
-0.97 0.03 43 68 22 76 8 0.90
~0.88 0.03 42 122 25 82 9 1.49
-0.77 0.03 . 43 90 21 58 6 1.55
-0.69 0.02 47 160 24 59 6 2.71
-0.62 0.02 40 180 26 57 6 3.15
-0.54 0.02 42 156 23 52 6 2.99
-O.44 0.03 40 180 21 48 5 3.74
~0.28 0.11 52 183 15 43 4 4.22

 

 

 

 

 

 

 

 

 

93

 

 

 

-2.0

 

 

 

-l.5

[Fol HI

-l.0

 

 

 

 

 

 

r1111d1111d11111111111111.1111G1111 fi 1 1 1 1 d 1 1 1 1 d 1 1 1 1 fi +11111111q1111d-14111d‘111
V
r 8 A v A
8 TI 8 TOII. A 8 IA
r v A r . A
.l I. m ' I. m ..I I
. .
r A v A v
r A 8 A r A
fi L v A f L
r A 8 _ _ _ _ u _ A
A v

.I l U I L u r l
. T.Cl. . . . II. . . . 70.2
v A m 8 m v A
v A V f u r

ll 8.. II 8. .
v A v

o e.
.. TI 1 a. .- IY. L a. u 8.0.. .-
v C a O A f n a C A r - a - A
r u I. A 8 .O- a n - A
r 80. . . 5. . T!
r II. A v 70.. A v i
I 70.. .. M .- 3. .- M .. II. .
v A r i A v A
v k v r
r A v A r
. ,0. . . .0. L . II.
IIII-IIIIDIIII-IIII-IIII-IIII-IIII M I I I P D I I I P P I I I M IPII-IIII-PIII-IIII-PI
. .2..2—...{83-8’
CO .8 a

0.0
Figure 9. Solution for Ur“. 010,, and their ratio as a function of abundance, using

the method of FW80 and the data of Table 4.

TABLE 5. FW80 Analysis Results.

94

 

 

 

([Fe/H]) Ogre/H] N um 08 (no. (disphoa I vrot/Utos I
(km 8'1) (km 8.1) (km 871) (km 8'1)

(1) (2) (3) (4) (5) (6) (7) (8)
-1.82 0.21 80 13 21 95 8 0.14
-1.28 0.13 82 46 21 100 8 0.46
-0.93 0.06 85 93 17 80 6 1.17
-0.73 0.05 90 121 16 60 4 2.02
—0.56 0.05 82 168 17 54 4 3.08
—0.35 0.12 92 182 12 45 3 4.02

 

 

 

 

 

 

 

 

 

95

1111.1:‘d1111“‘11.1111-1411‘1111 u

 

. TI .0
.I ll
.
v A
v A
v A
V
T I.”
r A.
r 10.. .
r A
. A0
I 1L
f... L.
Y
r A
i A
V
r __ A
..l 1M
r A
V A
.8! 8
f A

 

 

o ommomm
um... ..-

 

0.0

II II. .

31>

(Fe/Bl

 

v V V—V—

 

4
q
1
1

d
1
1
1
1

d
1
1
1
1

i

IIIIDIIII-III

l

A l A A

'15 -2.0

-l.0

 

 

o
5
l

w o
5

l

31.63.

[Fe/El

 

 

 

 

1111‘1‘11‘1111‘111‘1111 -
r
v
V A
V
I. 0
.. 1 1
.
v A
r A
v A
r A
x 1 ......
.
r A
. 1!.
v A
r
0.
r. I. In
v0. L .
V
v
. 70..
v ... A
4.
rl AL 0
.
8 .
v
. TOI.
DD’bL-b”-P-’b’b-;DI-FI’ M
5 4 3 1 1A 0
31.338\83n>

[Fe/HI

Figure 10. Solution for 2),“, 0'1”, and their ratio as a function of abundance, using

the method of FW80 and the data of Table 6.

96
seem that at least 40 objects per bin is suficient to obtain a value for v.0, with an

uncertainty that is close to the external errors in the velocities for the stars.

A potential weakness in the result presented here is the division in metallicty
that is used to separate the FHB and A stars. To explore the effect that this cutoff
has in the response of the rotation determination, the calculation was repeated
twice, using a division of [Fe/H] = —0.50 and [Fe/H] = -1.00 to separate the
FHB and A stars. A comparison of the results obtained is given in Figure 11,
and it can be seen that the effect on v.0; is to increase the range in [Fe/H] over
which the transition from halo to disk kinematics occurs. It is therefore difficult to
assess the sharpness of the transition, as it depends heavily on the division between
the two sub-samples. A better knowledge of the metallicity distribution of the A
type stars would establish this upper limit in metallicity for which the kinematics
become those of the ”extended disk.” The lower limit, i.e., the point at which the
transition begins, is clearer, as it is very unlikely that the A type stars are numerous
at metallicities this low. In general, it can be seen that, for the results in Figure
9, the transition occurs over the range —1.20 _<_ [Fe/ H] S —0.6, which is consistent

with the determinations of others.

As well, the values of v.0; over the range of metallicities in this determination
are consistent with those derived by others. For the bins which approach solar
metallicity, the result is consistent with ”normal” properties of the ”extended disk.”
This is noteworthy, as the stars which populate these high-metallicity bins are those

which are classified as A type. They seem to participate in the same systematic

97

 

 

 

 

““.“““‘“-““.““.““1“‘
L -
A. .
. TI0I..
r A
flu
.- .. z
r A c
. m
l
r A
.. 1 5
II
v i A .
T01 -
0.
... T: I l
t - a - L -
8 70.. .
. 70.. -
I
II .- 5
I 0
T0... . .
3..
IIIIDIPIIDIIII-IIII-IIIIDIIIIhIIII AN
AU
0 o o o o o o o
u o 5 o s s o
2 I. IA . Am.
dOHI>

lPe/Hl

 

 

'15 -2.0

-l.0

 

. 8

 

1‘11.1111‘1111-11‘1‘1‘11-1“1‘1‘1
r
.II
I
V
v
8
TI.
1
T
r TILWIL
r
Y
r
r TO...
. T0...
. 8.0..
8 TI
TI
I
v
V
Y
r 79..
IIIIDIIII-IIII-IIII.IIII-IIIIPII
0 o o m o o 0
fl 0 5 5 5
2 l I. .

31>

o

0

II.
.

[Fe/m

 

 

 

 

5
1111‘1111.1111‘1111‘111111111‘1111 %
r A
v TIIAVII. A

0.
I. ll 2
.
r
v
T0...
5
r 1 1h
. T01 ..
1
v A
v A
a.
1 TI 1 l
. II. .
. T01 .
. T0... .
. .+. .
T0... 5
II .8.
r
8.0..
IIID.’D......DrD'I-ID.’-DIDI-IDI o.
0
0 0 0 0 0 0 0 0
H 0 5 0 5 5 0
5‘ ll ll . Afl
.Oul>

lFo/Hl

—O.75

Figure 11. A comparison of derived values for v," for three different divisions
between the FHB and A type star samples: [Fe/H] = -0.5 (top), [Fe/H]

(middle) and [Fe/H] :2 —1.00 (bottom).

98
rotation that other stars have in this range of metallicity, which suggests that they
are normal constituents of this region of the galaxy. The implication of this is
that the ”extended disk” is a site of star formation, and is therefore not a fossil
remnant of the galaxy’s collapse. Recent findings on the density of gas within a few
kpc of the galactic plane would seem to allow for this conjecture (Robertson et al.,
1990). High velocity gas clouds with metallicities of up to [Fe/ H] = —O.5 have been

detected, using quasars and other extragalactic objects as probes.

Another worry in drawing conclusions from this sample is the question of
whether the sky is adequately covered over each of the metallicity bins. Figure
12 is a stripe histogram of the values of cos $.- for the bins of Table 4 (each stripe
on the histogram represents one star). Since this is the angle which projects the
line of sight velocity along the radial direction, it is also a measure of how much
weight to assign each bin in the determination of Ur“. It can be seen that there is
sufficient coverage in each of the bins to consider the derived rotational velocities
valid, although the general decrease in the coverage at higher metallicities stems

from the relative proximity of the A type stars to the sun.

 

99

«w. - -m «wu- - 411

“Hug “W .——1I|-l—LLLJLLIJU llllll

 

 

4's 4. a u u «I -u u a :3
u- a.
1” C ‘1.“ W0 ' '0."

 

um lrl L 11 "n Hllllllllll I IIII I llllllllllllllll
H

0M 4! 8 Cl In ~01 4. u t.
a o .
MID-1.40 ”-1.01

as '10 u u m on «A a“ u u
a. .
M881," ”-4.“

 

 

II II "H“ I” lllfll I III Ll _Jmn ll". L IIIL fi
f a 00.. '1! C II t.

4.. 'U I. II tl

MA...” .'1~D--cu

 

on «1| II I. «I «0 I.
u- c.-
WD - ~00 «Ire/III - 42!

t0 4.

Wu um i I .' ._ "MW
06. '0 I. ll '01 O. 0! o
C. '

Figure 12. Stripe histograms showing the range in cos 11:,- for each of the bins in
Table 4.

 

CHAPTER 7: SUMMARY AND CONCLUSIONS

It has been demonstrated that the spatial structure of the inner halo likely de-
parts from a random distribution of stars to distances of z 8 kpc. This conclusion
stems from probing the structure of the halo with the FHB I catalogue, although
it is unclear if the departure seen is a result of clumping of matter on small scales
(5 100pc), or perhaps the existence of remnants of disrupted groupings. The un-
certainty here lies in the coarse brightness (and therefore distance) determinations
for these stars at present, which would be remedied by an observing programme
aimed at acquiring photometry for these stars. The extension of the HK Survey to
the northern galactic hemisphere will prove beneficial in attacking this and other

problems in galactic studies.

The rotation of the halo system as traced by the FHB stars is consistent with de-
terminations which have been made for other populations of stars with the halo. The
A type stars which are identified in the HK Survey appear to be normal members
of the extended disk population, sharing the kinematics, distribution of metallicity
and location of other stars which have been identified as part of this component of

the galaxy. The transition in the halo from halo to extended disk kinematics

100

101
occurs from [Fe/H] z —1.2, and appears to extend to —0.80 S [Fe/H] _<_ -0.6, and
is smooth. This appears to confirm that the disk is decoupled from the halo, so
that the evolutionary picture of Eggen, Lyndon-Bell and Sandage would seem to be

at odds with this result.

LIST OF REFERENCES

LIST OF REFERENCES

Bahcall, J.N., and Soneira, RM. 1984. Ap.J.-S'uppl., 55, 67.
Beers, T.C., and Doinidis, SP. 1988. Bull-A.A.S., 19, 1030.

Beers, T.C., Preston, G.W., and Schectman, S.A., 1985 AJ. 90, 2089.
(BPS I).

Beers, T.C., Preston, G.W., and Schectman, S.A. 1990, in preparation.
(BPS II).

Beers, T.C., Preston, G.W., and Schectman, S.A. 1988, Ap.J.Suppl. 67
461. (FHB I).

Burstein, D., and Heiles, C. 1982. 11.7., 87, 1165.

Clube, S.V.M., and Watson, RC. 1979. M.N.R.A.S., 187, 863.

Eggen, O.J., Lyndon-Bell, D., and Sandage, A.R. 1962. Ap.J., 136, 748.
Henk, CS, and White, S.D.M. 1980, M.N.R.A.S. 193 295. (FW80).
Gilmore, G., Wyse, R.F.G., and Kuijken, K. 1989 Preprint-

Green, R.F. 1980. Ap.J., 238, 685.

Kilkenny, D. 1987 IAU Colloquium 95: The Second Conference on Faint
Blue Stars, A.G.D. Philip, D.S. Hayes and J .W- Liebert (eds.),
p.555, L. Davis Press.

Kinman, TD. 1959. M.N.R.A.S., 119, 559.

K00, DC, and Kron, R.G., 1982. Astron.Astrophys., 105, 107.

Koo, D.C., Kron, R.G., and Cudworth, KM. 1986. P.A.S.P., 98, 285.
Morrison, H.L., Flynn, C. and Freeman, KC. 1990, Proprint.

Noguchi, T., Maehara, H. and Kondo, M. 1980. Ann., Tokyo, Astron..
Obs. 18 55

Norris, J .E. 1986 Preprint-
Norris, J.R., and Ryan, SC. 1989. Ap.J., 340, 739.
Pier, J.R. 1982. A.J., 87, 1515.

102

103

Pier, J.R. 1983. Ap.J.Suppl., 53, 791.
Preston, G.W., Scectman, S.A. and Beers, T.C. 1990, Preprint-

Robertson, J.C., Schwarz, U.G., van Woerden, H., Murray, J .D., Mor-
ton, DC. and Hulsbosch, A.N.M. 1990 A.A.O. Preprint 252 (1990

November) .

Slettebak, A., and Stock, J. 1959, ”A Finding List of Stars of Spectral
Type F2 and Earlier in a North Galactic Pole Region,” Hamburg-
Bergedorf Obs. Publ-

Slettebak, A. and Brundage, R.K. 1971. A.J., 76, 338.
Sommer-Larsen, J., and Christensen, RR. 1987. M.N.R.A.S., 219, 537.
Sommer-Larsen, J., and Christensen, RR. 1989. M.N.R.A.S., 239, 441.
Thomas, P. 1989, M.N.R.A.S., 238 1319. (T89).

Woltjer, L. 1975. Astron.Astrophy.9-, 42, 109.

HICHIG

(in)