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Fundamental Questions in Astrophysics: Guidelines for Future UV Observatories

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Abstract

Progress of modern astrophysics requires the access to the electromagnetic spectrum in the broadest energy range. The Ultraviolet is a fundamental energy domain since it is one of the most powerful tool to study plasmas at temperatures in the 3,000–300,000 K range as well as electronic transitions of the most abundant molecules in the Universe. Moreover, the UV radiation field is a powerful astrochemical and photoionizing agent. The objective of this book is to describe the crucial issues that require access to the UV range.
FUNDAMENTAL QUESTIONS IN ASTROPHYSICS:
GUIDELINES FOR FUTURE UV OBSERVATORIES
Edited by:
ANAI.G´
OMEZ DE CASTRO and WILLEM WAMSTEKER
Reprinted from Astrophysics and Space Science
Volume 303, Nos. 1–4, 2006
Library of Congress Cataloging-in-Publication Data is available
ISBN 1-4020-4838-6 (hardbook)
ISBN 1-4020-4839-4 (eBook)
ISBN 978-1-4020-4838-6 (hardbook)
ISBN 978-1-4020-4839-4 (eBook)
Published by Springer,
P.O. Box 17, 3300 AA Dordrecht, The Netherlands.
Picture left: inserted as tribute to the late Willem Wamsteker who liked this image very much
Pictures right:
Top: Galaxy halo if the Universe reionized at redshift 15 or 6, by Kenji Bekki & Masashi Chiba, Tohuku University, Japan
Below: Sonic Point model of KiloHertz QPOs, by M. Colleman, F.K. Lamb & D. Psaltis
Below: Simulations of accretion disks, by J.F. Hawley, S.A. Balbus, J.M. Stone
Below: Simulations of the interaction of the accretion disk and the magnetized star inaTTauri System, by Brigitta von
Rekowsky & Axel Branderger
Bottom: Artist illustration of the evaporation of an exoplanet atmosphere, by European Space Agency and Alfred
Vidal-Madjar (Institut d’Astrophysique de Paris)
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Printed in the Netherlands
TABLE OF CONTENTS
Foreword 1–2
M.A. Barstow and K. Werner / Structure and Evolution of White Dwarfs and their Interaction with the Local
Interstellar Medium 3–16
Isabella Pagano, Thomas R. Ayres, Alessandro C. Lanzafame, Jeffrey L. Linsky, Benjam´ın Montesinos and
Marcello Rodon`o / Key Problems in Cool-Star Astrophysics 17–31
Ana I. G´omez de Castro, Alain Lecavelier, Miguel D’Avillez, Jeffrey L. Linsky and Jos´e Cernicharo / UV
Capabilities to Probe the Formation of Planetary Systems: From the ISM to Planets 33–52
Boris T. G¨aansicke, Domitilla de Martino, Thomas R. Marsh, Carole A. Haswell, Christian Knigge, Knox S.
Long and Steven N. Shore / Ultraviolet Studies of Interacting Binaries 53–68
Willem Wamsteker, Jason X. Prochaska, Luciana Bianchi, Dieter Reimers, Nino Panagia, Andrew C. Fabian,
Claes Fransson, Boris M. Shustov, Patrick Petitjean, Phillipp Richter and Eduardo Battaner / The Need
for Ultraviolet to Understand the Chemical Evolution of the Universe, and Cosmology 69–84
Rosa M. Gonz´alez Delgado / Starbursts at Space Ultraviolet Wavelengths 85–102
Noah Brosch, John Davies, Michel C. Festou and Jean-Claude G´erard / A View to the Future: Ultraviolet Studies
of the Solar System 103–122
Wolfram Kollatschny and Wang Ting-Gui / Active Galaxies in the UV 123–132
Ana I. G´omez de Castro, Willem Wamsteker, Martin Barstow, Noah Brosch, Norbert Kappelmann, Wolfram
Kollatschny, Domitilla de Martino, Isabella Pagano, Alain Lecavelier des ´
Etangs, David Ehenreich,
Dieter Reimers, Rosa Gonz´alez Delgado, Francisco Najarro and Jeff Linsky / Fundamental Problems in
Astrophysics 133–145
Norbert Kappelmann and J¨urgen Barnstedt / Guidelines for Future UV Observatories 147–151
F. Najarro, A. Herrero and E. Verdugo / Massive stars in the UV 153–170
Astrophys Space Sci (2006) 303:1–2
DOI 10.1007/s10509-006-9175-z
Foreword
C
Springer Science +Business Media B.V. 2006
Modern astrophysics is a mature science that has evolved
from its early phase of discovery and classification to a
physics-oriented discipline focussed in finding answers to
fundamental problems ranging from cosmology to the origin
and diversity of life-sustainable systems in the Universe. For
this very reason, progress of modern astrophysics requires
the access to the electromagnetic spectrum in the broadest
energy range. The Ultraviolet is a fundamental energy do-
main since it is one of the most powerful tool to study plas-
mas at temperatures in the 3,000–300,000 K range as well
as electronic transitions of the most abundant molecules in
the Universe. Moreover, the UV radiation field is a powerful
astrochemical and photoionizing agent.
The impact of UV instruments in astronomical research
can be clearly traced through the considerable success of
the International Ultraviolet Explorer (IUE) observatory and
successor instruments such as the GHRS and STIS spectro-
graphs on-board the Hubble Space Telescope (HST), or the
FUSE satellite operating in the far UV (90–120 nm range).
Of particular importance has been access to high resolution
R40,000–100,000 spectra providing an ability to study
the dynamics of hot plasma and separate multiple galactic,
stellar or interstellar spectral lines. Furthermore, the GALEX
satellite is providing new exciting views of UV sources.
This book describes the fundamental problems in modern
astrophysics that cannot progress without easy and wide-
spread access to modern UV instrumentation. Three among
them should be stressed by its relevance:
1. Extrasolar planetary atmospheres and astrochemistry in
the presence of strong UV radiation fields.
2. Chemical evolution of the Universe and the diffuse bary-
onic content.
3. The physics of accretion and outflow: the astronomical
engines.
The volume contains a series of review articles that analyze
the scientific requirements for modern UV instrumentation.
The first article summarizes the science case for UV astron-
omy. After, several articles targeting the major research fields
of astrophysics from Solar System to cosmological research
are included. These articles analyze why and which UV in-
strumentation is required to make progress in each field. The
book ends with a summary of the UV instrumentation de-
manded by the community and a brief update on techno-
logical requirements. All articles in this volume have gone
through the peer-review system of the journal “Astrophysics
and Space Science.”
This book contains the thoughts and work of the Network
for UV astrophysics (NUVA) and the UV community at-
large. By the end of 2002, a group of European astronomers
coming from a broad range of areas: from fundamental as-
trophysical research to observational expertise in the optical,
UV and X-ray ranges, as well as space instruments develop-
ment teams, joined efforts to evaluate the need to develop new
UV instrumentation for the coming decade in order to achieve
some of the major scientific objectives of the astronomical
community and make full profit of the large astronomical
facilities planned for other spectral ranges; this group set the
seed for the Network for UltraViolet Astrophysics (NUVA).
At that time, STIS was working in HST and the Cosmic Ori-
gins Spectrograph (COS) was being built for replacement;
HST was thought to last till 2010/12, FUSE was working
nominally and GALEX was close to be launched. However,
the big space agencies had no plans for new UV spectroscopic
missions unless the large optical/UV telescope included in
the NASA Origins plan targeted to enter development phase
in 2015–2020 and to be launched in the second quarter of this
century. Unfortunately, on Friday 16 January 2004 NASA
informed that the planned shuttle mission to service and up-
grade HST (SM-4) including the substitution of STIS by
Springer
2Astrophys Space Sci (2006) 303:1–2
COS, had been cancelled. Just few months later STIS failed.
No access to high resolution UV spectroscopy has been fea-
sible during most of 2005 since FUSE resumed observations
in November 2005 after the failure of the third of the four
onboard reaction wheels in December 2004. The NUVA was
officially established in January 2004. A key objective of the
NUVA is to run an exploratory analysis to define the sci-
entific requirements for future UV observatories. This book
contains the first outcome of this work and its publication
has been funded by the NUVA as a part of its activities. The
NUVA is within the Optical Infrared Co-ordination Network
for Astronomy (OPTICON) funded by the European Com-
missions 6th Framework programme under contract number
RII3-Ct-2004-001566.
In these two years, some respected colleagues and dear
friends who actively promoted and enthusiastically colabo-
rated in this project have passed away.
Willem Wamsteker was the director of ESA’s IUE Ob-
servatory until the mission terminated in 1996. After his di-
rection, the IUE data archive became the first fully internet
driven astronomical archive. Willem was a promotor of the
NUVA and a key initiator of the World Space Observatory
UV (WSO/UV) project; a 1.7 m UV telescope equipped with
state of the art instrumentation providing a factor of 10 im-
provement on the high resolution spectroscopic capabilities
of HST/STIS. WSO/UV will be launched in 2010 becoming
the only UV spectroscopic facility available world-wide. The
project is driven by a broad international colaboration led
by Russia (ROSCOSMOS); there is a significant European
participation in the project. Willem Wamsteker was also co-
editor of “Astrophysics and Space Science” and deeply in-
volved in the edition of this volume.
Marcello Rodon´o was director of the Obsservatorio As-
trofisico di Catania and a very relevant member of the “cool
stars” community. He was keen off multi-wavelength studies
and an active promotor of the WSO/UV project.
Michel Festou was director of l’Observatoire de Besan¸con
and a passionate researcher. His work on UV spectroscopy
of comets is reknowned world-wide. He actively colabo-
rated with the NUVA in the identification of the UV facilities
needed to make progress on Solar System Research.
We would like to dedicate this book to their memory.
Finally, we would like to thank our OPTICON colleagues
and very especially John Davies (OPTICON Project Scien-
tist) and Gerry Gilmore (OPTICON P.I.) for their support.
We also acknowledge the support of the Universidad Com-
plutense de Madrid which is hosting and maintaining the
NUVA site (www.ucm.es/info/NUVA).
Prof. Ana In´es G´omez de Castro
NUVA, chair
Madrid, March 1st, 2006
Springer
Astrophys Space Sci (2006) 303:3–16
DOI 10.1007/s10509-006-2061-x
ORIGINAL ARTICLE
Structure and Evolution of White Dwarfs and their Interaction
with the Local Interstellar Medium
M.A. Barstow ·K. Werner
Received: 11 March 2005 / Accepted: 11 August 2005
C
Springer Science +Business Media B.V. 2006
Abstract The development of far-UV astronomy has been
particularly important for the study of hot white dwarf stars.
A significant fraction of their emergent flux appears in the far-
UV and traces of elements heavier than hydrogen or helium
are, in general, only detected in this waveband or at shorter
wavelengths that are also only accessible from space. Al-
though white dwarfs have been studied in the far-UVthrough-
out the past 25 years, since the launch of IUE,onlyafew
tens of objects have been studied in great detail and a much
larger sample is required to gain a detailed understanding of
the evolution of hot white dwarfs and the physical processes
that determine their appearance. We review here the current
knowledge regarding hot white dwarfs and outline what work
needs to be carried out by future far-UV observatories.
Keywords Ultraviolet astronomy ·Space missions ·White
dwarfs
1. Introduction
1.1. White dwarf stars and the local interstellar medium
in astrophysics
Many white dwarfs are among the oldest stars in the Galaxy.
Their space and luminosity distributions help map out the
history of star formation in the Galaxy and can, in princi-
ple, determine the age of the disk by providing an important
M.A. Barstow ()
Department of Physics and Astronomy, University of Leicester,
University Road, Leicester LE1 7RH, UK
e-mail: mab@star.le.ac.uk
K. Werner
Universit¨at T ¨ubingen, FRG, T¨ubingen
lower limit to the age of the Universe. Recently, it has been
suggested that cool white dwarfs may account for a substan-
tial fraction of the missing mass in the galactic halo (Op-
penheimer et al., 2001). White dwarfs are also believed to
play a role in the production of type Ia supernovae, through
possible stellar mergers or mass transfer. Although no can-
didate precursor system has yet been found, they are used
as “standard candles” to measure the distances of the most
remote objects in the Universe and to imply a non-zero value
for the cosmological constant. It is intriguing that such lo-
cal, low luminosity objects as white dwarf stars play a key
role in some of the most important cosmological questions
of our day, concerning the nature and fate of the Universe.
To understand and calibrate cosmologically important as-
pects of white dwarfs, such as their cooling ages, masses and
compositions, we require a thorough understanding of how
their photospheric compositions evolve. Atmospheric metal
abundances affect cooling rates and bias the determination
of temperature and surface gravity. Reliable masses can only
be derived from accurate effective temperature and surface
gravity measurements. Importantly, metals are difficult to de-
tect in cool white dwarfs but still play an important role in
cooling. Therefore, abundances measured in hotter stars pro-
vide important data as to what species may be present and in
what quantities.
Interstellar gas is a fundamental component of the Milky
Way and other galaxies. The local interstellar medium
(LISM) is particularly important because it is close enough
to us for detailed examination of its structure. Indeed, it is the
only part of the Universe, apart from the solar system, that we
can study directly, from the penetration of neutral particles
through the heliopause. Study of the composition of the local
interstellar gas tells us about the evolution of the Universe and
our galaxy. While modified considerably since the formation
of the solar system, the processes that have shaped the ISM
Springer
4Astrophys Space Sci (2006) 303:3–16
must have also affected the solar system and possibly even
the evolution of life. For example, the local radiation field
and gas properties of the LISM set the boundary conditions
for the heliosphere. Therefore, it regulates the properties of
the interplanetary medium and the cosmic ray fluxes through
the solar system (see Frisch et al., 2002), which may affect
the terrestrial climate (Mueller et al., 2003).
The relationship between white dwarfs and the LISM is
intimate. The process of producing white dwarfs in the disk
substantially enriches the content of the local ISM (Mar-
gio, 2001) and white dwarfs contribute significantly to the
total cosmic abundance of metals (Pagel, 2002). They also
provide a significant fraction of the total flux of ionizing ra-
diation within the LISM and will affect the ionization state
of the interstellar gas, even if the LISM is not in ionization
equilibrium. At the same time, white dwarfs may be accret-
ing material from the denser clouds in the ISM, which will
modify their photospheric abundances. Finally, white dwarfs
are ideal background sources for direct study of the LISM.
1.2. Physical structure and classification of white
dwarfs
Astronomers have known about the existence of white dwarf
stars for 150 years, since the discovery of a companion to the
brightest star in the sky, Sirius. Studying the regular wave-
like proper motion of Sirius, Bessel revealed the presence
of a hidden companion, with the pair eventually resolved by
Clark in 1862 and the orbital period revealed to be 50 years.
Later, the first spectroscopic study of the system by Adams
revealed a great enigma. While the temperature of Sirius B
was found to be higher than Sirius A, it was apparently 1000
times less luminous. This result could only be explained if
Sirius B had a very small radius, 1/100 Rand similar to
that of the Earth. However, with a known mass of 1Mfor
Sirius B, the implied density of 1400 Tonnes m3was well
above that of any form of matter known to 19th and early
20th century physicists. Within a number of years, a handful
of stars similar to Sirius B were found, mostly in binaries.
The term “white dwarf” was coined later, on the basis of the
visual color and small size, when compared to other stars.
The answer to the riddle of their structure stems from the
development of the quantum statistical theory of the electron
gas by Enrico Fermi and Paul Dirac. When atoms in a mate-
rial are sufficiently close together, their most weakly bound
electrons move freely about the volume and can be consid-
ered to behave like a gas. Almost all the electrons in the gas
will occupy the lowest available energy levels. Any states
with the same energy are said to be “degenerate.” Hence, this
gas is known as the “degenerate electron gas.” Under condi-
tions of extreme pressure the electrons are forced to occupy
space much closer to the nuclei of the constituent atoms than
in normal matter, breaking down the quantised structure of
the energy levels. However, according to the Pauli Exclusion
Principle, no two electrons can occupy the same quantum
state, which has a finite volume in position-momentum space,
so there is a limit where a repulsive force arising from this, the
“electron degeneracy pressure”, resists further compression
of the material. Fowler (1926) showed that this pressure could
support a stellar mass against gravitational collapse and pro-
posed that this might explain the existence of white dwarfs.
Combining this insight with the equations of stellar structure,
Chandrasekhar (1931, 1935) determined the mass–radius re-
lation for white dwarfs and the maximum mass (1.4 M) that
electron degeneracy pressure could support against gravity,
the Chandrasekhar limit. As this Nobel Prize winning work
has been extended by subsequent developments, the basic
ideas remain unchanged. However, importantly, they have
also hardly been tested by direct observation.
Theoretical and observational study of stellar evolution
has placed white dwarfs as one possible end point of the pro-
cess. In general terms, all stars with masses below about 8
times that of the Sun will pass through one or more red giant
phases before losing most of their original mass to form a
planetary nebula. The remnant object, a white dwarf, is the
core of the progenitor star. In the absence of any internal
source of energy, the temperature of a white dwarf, after its
birth, is determined by how rapidly stored heat is radiated
into space. Estimates of white dwarf cooling times indicate
that it will take several billion years for the stars to fade to in-
visibility. Hence, white dwarfs are among the oldest objects
in the galaxy. Since, the galaxy is younger than the cool-
ing timescales; the lowest temperature (oldest) white dwarfs
yield a lower limit to its age.
The basis for understanding the nature of most stars is
analysis of their optical spectra and classification according
to the characteristics revealed. A number of physical pro-
cesses can alter the atmospheric composition of a white dwarf
as it cools. As noted by Schatzmann (1958), the strong grav-
itational field (log g8 at the surface) causes rapid down-
ward diffusion of elements heavier than the principal H or He
component. Hence, Schatzmann predicted that white dwarf
atmospheres should be extremely pure. Consequently, the
spectra should be devoid of most elements, showing signa-
tures of only hydrogen and, possibly, helium. White dwarfs
are thus divided into two main groups according to whether
or not their spectra are dominated by one or other of these
elements. The hydrogen-rich stars are given the classifica-
tion DA, while the helium-rich white dwarfs are designated
DO if HeII features are present (hotter than about 45,000 K)
and DB if only HeI lines are visible. Small numbers of hy-
brid stars exist, with both hydrogen and helium present.
In these cases, two classification letters are used, with the
first indicating the dominant species. For example, DAO
white dwarfs are mostly hydrogen but exhibit weak HeII
features.
Springer
Astrophys Space Sci (2006) 303:3–16 5
Fig. 1 Schematic description of
the production of H-rich and
He-rich branches of white dwarf
evolution
The above classification scheme applies only to white
dwarfs with temperatures above 10,000 K, when the H
and He energy levels are sufficiently populated above the
ground state to produce detectable features at visual wave-
lengths. At lower temperatures, although H and/or He may
be present these elements are no longer directly detectable.
Cool white dwarfs are divided into three main groups. DC
white dwarfs have continuous, completely featureless spec-
tra; DZ white dwarfs are He-rich and have only metal lines
visible; DQ white dwarfs show carbon features. Figure 1
summarises the current thoughts regarding the relationships
between the hot white dwarf groups along with the princi-
pal mechanisms that provide evolutionary routes between
them.
Gravity is not necessarily the only influence on photo-
spheric composition. It can be countered by radiation pres-
sure which acts outward to support heavy elements in the
atmosphere, a process termed “radiative levitation.” Another
mechanism that can mix elements that have settled out in
the stellar atmosphere is convection. If the convective zone
reaches down to the base of the atmosphere then heavy el-
ements can be dredged back up into the outer atmosphere.
A further complication is that material can also be accreted
from the interstellar medium.
1.3. The structure of the local interstellar medium
The Sun is embedded in the Local Interstellar Cloud (LIC), a
warm (T10,000 K), low-density (n0.1cm
3) region,
which is observed in projection toward most, but not all,
nearby stars. Models of the LIC (Redfield and Linsky, 2000)
show the Sun located just inside the edge of the LIC in the di-
rection of the Galactic Center and toward the North Galactic
Pole (NGP). The LIC is surrounded by a hot 106K substrate
and appears to be part of the expanding cloud complex rep-
resenting the Loop I radio shell, which is either a supernova
remnant or a wind blown shell from the Sco-Cen Associa-
tion (Fig. 2). Indeed, evidence for shocks in the LIC within
the past 2 Myr is supported by EUV observations of ion-
ized interstellar HeII (Lyu and Bruhweiler 1996; Barstow
et al., 1997). Also, surrounding the LIC is the low den-
sity Local Bubble with dimensions of >200 pc. Soft X-ray
and OVI observations imply this region is largely filled with
a hot (T106K), extremely low-density (n102cm3)
plasma. Frisch (1995) and Bruhweiler (1996) discuss possi-
ble origins of Loop I and the Local Bubble, but we need to
resolve two key questions. What is the geometry of the Lo-
cal Bubble? What are the distribution, physical conditions,
and kinematics of clouds in the Local Bubble? The answers
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6Astrophys Space Sci (2006) 303:3–16
Fig. 2 View of the LISM within
400 pc of the Sun. The Sun is
located at the end of the dashed
line extending from the star β
CMa. The dashed line is
200 pc in length. The region of
the Local Void or Local Bubble
is denoted. Most of the volume
of the LISM is composed of
very low-density gas. The radio
loop structures are further
indicated. The stars of Sco-Cen
are interior to Loop I. Regions
of high dust extinction, regions
where dense clouds are found,
are also marked. From
Bruhweiler (1996)
to these questions also bear directly on whether interstellar
accretion is possible in specific white dwarfs.
2. The complex nature of the DA white dwarf
population
2.1. Composition and structure of the stellar
photospheres from UV spectroscopy
The emergence from the Asymptotic Giant Branch (AGB) of
the two H- and He-rich groups, is qualitatively understood
to arise from differences in precise evolutionary paths. The
He-rich group undergoes a late helium shell flash, which
causes convective mixing of the outer envelope. Hydrogen
is ingested and burnt or diluted by this process (e.g., Herwig
et al., 1999). The majority of young white dwarfs are of DA
type, out numbering the He-rich objects by a ratio of 7:1
(Fleming et al., 1986; Liebert et al., 2005). The more recent,
deeper, Sloan Digital Sky Survey appears to yield an even
higher value of 9:1 for the ratio of DAs to the total number
of DO and DB white dwarfs (Kleinman et al., 2004). How-
ever, the complex relationships between these groups and
a demonstrable temperature gap in the He-rich cooling se-
quence cannot yet be readily explained. For example, the ratio
of H-rich to He-rich progenitors (4:1) is lower than that of
the white dwarfs into which they evolve (Napiwotzki, 1999).
Furthermore, a gap in the He-rich sequence, between 45000
and 30000 K, separating the DO and DB white dwarfs (We-
semael et al., 1985; Liebert et al., 1986; Dreizler and Werner,
1996) implies that He-rich white dwarfs must temporarily be
seen as DA stars due to some physical process.
It is well established that the hottest (T50,000 K)
white dwarfs possess significant abundances of elements
heavier than He in their atmospheres. Typical ranges are
shown in Table I for the most important elements. Mass frac-
tions are given, with respect to H for DAs and He for DOs,
since they are essentially independent of nuclear burning,
if the element does not participate in the process. Observed
abundances may vary considerably with temperature, as dis-
cussed later. Hence, we list typical values and comment on
the possible ranges in each white dwarf pathway. Studying
how photospheric abundance patterns change as these stars
cool can delineate the evolutionary pathways followed by
white dwarfs. For example, abundance enhancements in C,
N and O may be a sign of near exposure of the C O core
during the AGB phase or of a late He-shell burning episode.
Springer
Astrophys Space Sci (2006) 303:3–16 7
Table 1 Abundances of
important heavy elements
observed in DA and DO white
dwarfs
DA DO
Element Typical value Comment Typical value Comment
C10
6to 105Zero in coolest 3 ×103
N10
6From zero to 1025×105In a few cases
at extremes
O10
4Zero in coolest 105
Si 104Seen at most temps 104
down to 108
Fe 2–5 ×104Zero in coolest Not detected in any DOs
Ni 1–2 ×105Zero in coolest 104
Note that the values are much
more uncertain for the less
well-studied DO stars
Fig. 3 1230 ˚
A to 1280 ˚
A region
of the STIS spectrum of
REJ0558-373, showing
photospheric absorption lines of
N V (1238.821/1242.804 ˚
A) and
large numbers of Ni lines. The
best-fit synthetic spectrum is
shown offset for clarity. The
strong line near 1260 ˚
A, present
in the observation but not in the
model, is interstellar Si II (from
Barstow et al., 2003b)
Focused spectroscopic studies of the hottest white dwarfs
have begun to establish the general pattern of abundance with
respect to evolutionary status. Ultraviolet observations have
played an essential role by providing access to absorption
lines of elements heavier than H or He. Such features are
not present in optical or IR data except where photospheric
abundances are unusually high. Indeed, He is also hard to de-
tect, requiring abundances in excess of a few times 103in
the visible band. Typical detection limits in the UV are two
orders of magnitude lower. Therefore, the most important
and useful transitions, particularly many resonance lines of
elements heavier than H and He, lie in the far-ultraviolet (far-
UV, 1000–2000 ˚
A). However, since the lines are expected
to be weak and narrow, they are only normally visible at
high resolution (R>20,000). An example is a small sec-
tion of the high-resolution spectrum of the DA white dwarf
REJ0558-373, recorded with the Space Telescope Imaging
Spectrograph (STIS) onboard HST (Fig. 3), which shows the
interstellar 1260.4 ˚
A line of Si II together with photospheric
NV.
A survey of 25 hot DA white dwarfs, based on IUE and
HST data (Barstow et al., 2003b), shows that, while the
presence or absence of heavy elements largely reflects what
would be expected if radiative levitation were the support-
ing mechanism, the measured abundances do not match pre-
dicted values very well. These and earlier results are forcing
us to confront complexities in the real physical structure of
the stars. For example, it has become clear that the shape and
strength of the Balmer line profiles, from which Teff and log g
(and, indirectly, mass) are measured, are dependent on the
stellar photospheric abundances, requiring a self-consistent
analysis of each individual star based on data acquired at all
wavelengths (Barstow et al., 1998). Furthermore, we now
have direct observational evidence (Barstow et al., 1999;
Dreizler and Wolff, 1999) that photospheric heavy elements
are not necessarily homogeneously distributed (by depth) and
that more complex stratified structures must be considered.
While almost all stars hotter than 50,000 K contain heavy
elements, as expected, there is an unexplained dichotomy at
lower temperatures, with some stars having apparently pure
H atmospheres and others detectable quantities of heavy el-
ements (Barstow et al., 2003b, e.g., Fig. 4). In many of these
objects the photospheric opacity does not reveal itself in EUV
photometric or spectroscopic observations, implying that the
observed material resides in a thin layer in the uppermost re-
gion of the photosphere (see Holberg et al., 1999; Barstow et
al., 2003b). The effect of this stratification can be observed di-
rectly with high resolution UV spectroscopy. Figure 5 shows
Springer
8Astrophys Space Sci (2006) 303:3–16
Fig. 4 Measured abundances of
nitrogen (number ratio with
respect to hydrogen) as a
function of Teff for a sample of
25 DA white dwarfs (from
Barstow et al., 2003b)
Fig. 5 The STIS E140M
spectrum of REJ1032+532 in
the region of the NV resonance
doublet (histogram), compared
to the predicted line profiles
from a homogeneous model
atmosphere calculation (from
Holberg et al., 1999)
the HST/STIS spectrum of the white dwarf REJ1032+532,
compared to that predicted by a homogeneous model stel-
lar atmosphere. The N V line profiles in the model have
significantly broader wings than are observed, while a strat-
ified model (not shown) gives a much better match to the
data.
Recently, important progress has been made in incorporat-
ing radiative levitation and diffusion self-consistently into the
atmosphere calculations (Dreizler and Wolff, 1999; Schuh et
al., 2002). This work reconciles the overall spectral distri-
bution across the soft X-ray, EUV and far-UV bands with
the models (a problem with homogeneously distributed el-
ements) and explains the level of stratification inferred for
various elements. However, the abundance predictions do
not match observations for the known gravity of each star
observed, and agreement requires a higher surface grav-
ity than allowed by the optical data. In particular, we can-
not account for the large observed compositional differ-
ences between stars with identical temperature and surface
gravity.
In almost all the hot DA white dwarfs observed, the high
ionization resonance lines arising from photospheric heavy
elements have blue-shifted components, indicating that there
is some circumstellar gas present (Bannister et al., 2003).
Whether or not this material is a remnant of the planetary
nebula or due to ongoing mass-loss is unresolved. This has
important consequences for our basic understanding of stel-
lar composition. The effects of mass-loss, in the form of
weak winds ejecting material into the local ISM, and direct
accretion of material from the ISM are likely to be of great
importance in providing a plausible framework that can ex-
plain measured abundances.
Springer
Astrophys Space Sci (2006) 303:3–16 9
Hot He-rich DO white dwarfs are the progeny of stars
which constitute an interesting spectral class. Much of the
detailed physics involved in studying their atmospheres is
similar to that of the DA white dwarfs. However, the levi-
tation of heavy elements in these objects is far from under-
stood. The DO progenitors are the so-called PG1159 stars,
hot hydrogen-deficient objects, some of which represent the
hottest white dwarfs known, while others are still burning
helium in a shell.
Optical and UV spectral analyses have shown that the ef-
fective temperatures of PG1159 stars range between 75,000
and 200,000 K and the derived surface gravities are between
log g=5.5 and 8 (Werner, 2001). They are probably the out-
come of a late helium-shell flash, a phenomenon that drives
the fast evolutionary rates of three well-known objects (FG
Sge, Sakurai’s object, V605 Aql). Flash-induced envelope
mixing produces a H-deficient stellar surface (Herwig et al.,
1999). The He-shell flash transforms the star back to an AGB
star (born-again AGB star) and the subsequent, second post-
AGB evolution explains the existence of Wolf–Rayet cen-
tral stars of planetary nebulae (spectral type [WC]) and their
successors, the PG1159 stars. The photospheric composition
then essentially reflects that of the region between the H-
and He-burning shells in the precursor AGB star. It is dom-
inated by He, C, and O. Typical values are He =33%, C =
50%, O =17% (by mass), however, a considerable spread
of abundance patterns is observed, pointing to complicated
processes in the stellar interiors. Few stars show traces of
nitrogen (1%) or considerable amounts of residual hydrogen
(about 25%).
PG1159 stars provide the unique possibility of studying
the chemistry in the intershell region between the H- and
He-burning shells that is created after complicated and still
poorly understood burning and mixing processes during the
AGB phase. Usually the intershell material remains hidden
within the stellar interior. During the third dredge-up on the
AGB, however, intershell material can get mixed into the
convective surface layer and appears on the stellar surface,
though in rather diluted abundances. Nevertheless, this pro-
cess defines the role of AGB stars as contributors of nucle-
arly processed matter to the Galaxy. The motivation to study
PG1159 stars is based on the fact that these objects directly
display their intershell matter. It can be expected that grav-
itational settling is not affecting the composition, because
of ongoing mass-loss and convective motions. However, the
quantitative interpretation of the abundance analyses is still
premature because evolutionary calculations through a final
He-shell flash including a full nuclear network are not yet
available.
High-resolution UV spectroscopy was crucial in making
surprising discoveries which provide essential constraints to
calibrate theoretical modeling of stellar evolution. Generally,
UV spectra of PG1159 stars show only few photospheric
(absorption) lines, mainly from He II, C IV, O VI, and Ne
VII. Some of them display shallow N V lines and in many of
them we see sulfur. The S VI 933/944 ˚
A doublet in K1-16, for
example, suggests a solar abundance, which is in line with
the expectation that S is not affected by nuclear processes.
Silicon was also identified in some objects (Reiff et al., 2005),
but detailed abundance analyses remain to be done.
A very surprising result of UV spectroscopy was the
detection of a significant iron deficiency (1–2 dex) in the
three best studied PG1159 stars (Miksa et al., 2002). Obvi-
ously, iron was transformed to heavier elements in the in-
tershell region of the AGB star by n-captures from the neu-
tron source 13C(α,n)16 O (Herwig et al., 2003). Subsequently,
several other studies have also revealed an iron-deficiency
in [WC]-type central stars, which matches our picture that
these stars are immediate PG1159 star progenitors. Another
important result was accomplished by the identification of
one of the strongest absorption lines seen in FUSE spec-
tra of most PG1159 stars, located at 973.3 ˚
A. It is a Ne
VII line (Fig. 6) that allowed us to assess the neon abun-
dance in a large sample of objects (Werner et al., 2004).
Fig. 6 Discovery of a neon line (left panels) and a fluorine line
(right panels) in the hydrogen-deficient PG1159-type central star
PG1520+525 (top panels) and in the hydrogen-normal central star of
NGC 1360 (bottom panels). The neon and fluorine abundances in the
PG1159 star (given as mass fractions in the panels) are strongly en-
hanced, namely 20 times and 250 times solar, respectively, whereas
they are solar in NGC 1360. Note the strong Fe VII line (not included
in the models) at 1141.4 ˚
A in NGC 1360, which indicates a solar iron
abundance (Hoffmann et al., 2005). It is not detectable in the PG1159
star, probably due to a subsolar Fe abundance (Werner et al., 2004)
Springer
10 Astrophys Space Sci (2006) 303:3–16
It turns out that neon is strongly overabundant, (2%, i.e.,
20 times solar). This result clearly confirms the idea that
PG1159 stars indeed exhibit intershell matter. Neon is pro-
duced in the He-burning environment by two α-captures of
nitrogen, which itself resulted from previous CNO burning:
14N(α,γ)18 F(e+ν)18O(α,γ)22 Ne.
There are still many photospheric lines in UV spectra of
PG1159 stars which remain unidentified. Some of them may
stem from yet unknown Ne VII lines, or even of elements
which have not been detected in these stars at all. The latest
identification is that of a feature at 1139.5 ˚
A, which appears
rather strong in some objects. It is a line from highly ionized
fluorine (Fig. 6) and large overabundances (up to 250 times
solar) were derived for a number of PG1159 stars. This line
was also identified in “normal” hydrogen-dominated cen-
tral stars and, in contrast, solar fluorine abundances were
found (Werner et al., 2005). This again is a clear proof that
we see intershell matter on PG1159 stars. According to re-
cent calculations by Lugaro et al. (2004), their stellar models
show an effective fluorine production and storage in the in-
tershell, leading to abundances that are comparable to the
observed PG1159 abundances of fluorine. The general prob-
lem for fluorine production is that 19F, the only stable F
isotope, is rather fragile and readily destroyed in hot stel-
lar interiors by hydrogen via 19F(p,α)16 O and helium via
19F(α,p)22 Ne. The nucleosynthesis path for F production in
He-burning environments of AGB and Wolf–Rayet stars is
14N(α,γ)18 F(β+)18O(p,α)15 N(α,γ)19F.
All this underlines that AGB stars which dredge up ma-
terial from the intershell are contributing to the Galactic flu-
orine content (together with Wolf–Rayet stars and type II
SNe). This is completely in line with the detected fluorine
overabundances (up to 30 times solar) found from IR spectra
in AGB stars (Jorissen et al., 1992). To what extent PG1159
stars themselves return F to the ISM remains to be estimated.
The life time of a born-again AGB star is short in compari-
son to a usual AGB star, however, the fluorine fraction in the
mass lost by a wind of the former is much higher.
2.2. White dwarf masses and radii and the role of UV
imaging observations
Two of the most important physical parameters that can
be measured for any star are the mass and radius. They
determine the surface gravity by the relation g=GM/R2.
Hence, if log gis measured the mass can be calculated pro-
vided the stellar radius is known. One outcome of Chan-
drasekhar’s original work on the structure of white dwarfs
was the relationship between mass and radius, arising from
the physical properties of degenerate matter. Further theo-
retical work yielded the Hamada–Salpeter zero-temperature
mass-radius relation (Hamada and Salpeter 1961). How-
ever, white dwarfs do not have zero temperature, indeed
many are very hot. Hence, the Hamada–Salpeter relation
is only a limiting case and the effects of finite tempera-
ture need to be taken into account. Evolutionary calcula-
tions, where the radius of a white dwarf of given mass de-
creases as the star cools, have been carried out by Wood
(1992, 1995), Bl¨ocker (1995), Bl¨ocker et al. (1997) and
others.
Measurements of the surface gravity of samples of white
dwarfs show that the distribution of log gvalues and, there-
fore, of mass is very narrow (e.g., Bergeron et al., 1992;
Napiwotzki et al., 1999), with a peak mass of 0.6 M. This
is a direct consequence of the evolution of single stars, with
masses from 1 Mup to 8M. While the details of the re-
lationship between the initial mass of the progenitor star and
the final white dwarf mass are not particularly well under-
stood, it is clear that the small dispersion in the white dwarf
masses is related to a similarly small range of stellar core
masses and the fact that most of the outer stellar envelope
is expelled through several phases of mass loss along the
AGB. Importantly, any white dwarf with a mass outside the
approximate range 0.4–1.0 Mcannot arise from single star
evolution and must have an origin in a binary, where mass
exchange has taken place.
The basic model of the white dwarf mass-radius relation
is often used to derive masses from the spectroscopic mea-
surements of effective temperature and surface gravity (e.g.,
Bergeron et al., 1992; Napiwotzki et al., 1999). While this
is not in serious doubt, opportunities for direct observational
tests of the work are rare. This is particularly true of models
that take into account the finite stellar temperature and details
of the core/envelope structure, discussed above. Varying the
assumed input parameters in these models can lead to quite
subtle, but important differences in the model predictions.
To test these requires independent measurements of white
dwarf mass, which can be compared with the spectroscopic
results. Such information can be obtained spectroscopically
by measuring the gravitational redshift of absorption lines in
the white dwarf atmosphere (Vgr[km s1]=0.636 M/R,M
and Rin Solar units), but this is only possible if the systemic
radial velocity is known as a reference point. Generally, this
is only the case if the white dwarf is part of a binary system
and then there is also the possibility of obtaining independent
dynamical information on the white dwarf mass from the sys-
tem orbital parameters. An additional important constraint is
knowledge of the stellar distance.
The four best white dwarf mass determinations, where we
have the most complete and accurate information, are for 40
Eri B, Procyon B, V471 Tauri B and Sirius B, where we can
combine the assembled data with the Hipparcos parallax to
test the mass radius relation (Fig. 7). While there is good
agreement between the observation and theory, there nev-
ertheless remains a high degree of uncertainty in the mass
determinations. As a result, for example, it is not possible
Springer
Astrophys Space Sci (2006) 303:3–16 11
Fig. 7 Comparison of mass
estimates for 40 Eri B, Procyon
B, V471 Tau B and Sirius B
with the evolutionary models of
Wood (1995), displayed at
various temperatures and with
“thin” and “thick” H envelopes.
The solid limiting curve
represents the Hamada–Salpeter
zero temperature relation for a
carbon core (figure produced by
Jay Holberg)
Fig. 8 Wide Field Planetary
Camera image of the binary 56
Per, where each component (A
& B) is itself resolved into a pair
(left). (right) Successive images
of the Aa/Ab pair taken 18
months apart clearly show the
orbital motion of the system
to distinguish between different models, such as those with
“thin” or “thick” H envelopes.
We have such complete data (dynamical masses, gravita-
tional redshifts and accurate parallaxes) for only a very few
white dwarfs. Therefore, it is important to extend the sample
of white dwarfs for which we have to more objects and, pos-
sibly, explore a wider range of masses and temperatures. The
role of direct imaging in the UV as means of discovering
new systems is particularly important. For example, a ma-
jor result of the EUV sky surveys was the discovery of many
unresolved binary systems containing white dwarfs and com-
panion spectral types ranging from A to K (e.g., Barstow et
al., 1994; Burleigh et al., 1997; Vennes et al., 1998). In visi-
ble light the presence of a companion of spectral type earlier
than mid-K will swamp the signature of the white dwarf
making it undetectable. In the EUV, where the companion
flux is generally negligible, the white dwarf stands out very
clearly. The UV wavelength range is an even more efficient
way of searching for these binaries, as the interstellar opacity
is much lower than in the EUV, and the GALEX sky survey
is finding many examples.
The large difference in the visual magnitudes makes these
systems generally impossible to resolve with ground-based
observations. However, in the UV where the contrast is far
better it is possible to measure their separations, or at least
provide improved constraints. HST Wide Field Planetary
Camera 2 images of 18 binary systems have resolved 9 ob-
jects (Barstow et al., 2001a). Figure 8 shows one of the most
interesting examples, 56 Per, a known binary in which each
component has been resolved into a pair, making it a quadru-
ple star system. The white dwarf is a companion to 56 Per
A and is labeled 56 Per Ab in the image. At a distance of
42 pc, the measured 0.39 arcsec separation indicates a binary
period of 50 years for the Aa/Ab system. Therefore, the or-
bital motion of the two stars should be readily apparent with
repeated exposures on timescales 1–2 years, from which
a dynamical white dwarf mass can ultimately be obtained.
This is clearly demonstrated in Fig. 8 which shows a zoomed
view of the Aa/Ab pair from the main image and, on a similar
scale, a second image obtained 18 months after the first.
Continued monitoring of systems like 56 Per will eventually
yield the orbital parameters and dynamical determinations
Springer
12 Astrophys Space Sci (2006) 303:3–16
of the component masses. This requires continued access to
UV imaging.
3. Problems of white dwarf evolution
Clearly significant progress has been made in the study of
white dwarf stars through ultraviolet observations. However,
these have not yet given us access to a detailed understand-
ing of the important physics because they have been limited
in scope. First the total number of objects studied with the
necessary spectral resolution and signal-to-noise is small.
Therefore, we do not know how representative of the general
population individual objects or small groups of objects may
be. This is exacerbated by the fact that there are strong selec-
tion effects present within the existing samples. For example,
most hot DA white dwarfs have been observed because of ex-
pectations that significant quantities of heavy elements were
present in their photospheres, and, as a result are mostly stars
with temperatures above 50,000 K. Therefore, a number of
significant questions regarding the evolution of white dwarfs
remain to be solved:
rWhat is the origin of the DO–DB gap and the relationship
between the H- and He-rich white dwarf branches? For
example, do DO white dwarfs appear as DAs through float
up of residual H?
rWhat mechanisms determine the compositions of the DA
white dwarfs as they cool? Radiative levitation and gravi-
tational diffusion are clearly important, but why do some
stars of the same apparent temperature and gravity have
widely differing compositions? Is accretion (from the ISM
or a companion) an important mechanism? Do the abun-
dance differences reflect differences in progenitor compo-
sition/prior evolution of the progenitor?
rWhat are the metal abundances in the He-rich white
dwarfs? Are the PG1159 stars and DOs part of a single
sequence or do they represent the separate progenitor evo-
lutionary paths. Do all hot He-rich white dwarfs eventually
become DBs?
rWhat is the 3-D structure of the local ISM? How do white
dwarfs interact with it and exchange material. Does mass-
loss continue beyond the planetary nebula phase? Is appar-
ently circumstellar material detected in some of the hottest
white dwarfs evidence of such mass-loss or merely the
residual signature of PN material?
rWhat is the initial-final mass relation for both H- and He-
rich white dwarfs and what is the upper limit on the possible
progenitor mass? Does the theoretical mass-radius relation
for white dwarfs stand up to close scrutiny? Do the most
massive white dwarfs have exotic core compositions be-
yond the C/O product of helium burning?
All white dwarfs that have ever been studied in the UV re-
side within our own galaxy and must have emerged from
stellar populations with different ages and environments. To
solve the outstanding problem and make significant further
progress in the study of white dwarfs requires a substan-
tial enlargement of the sample, to properly examine the full
range of temperatures, gravities and possible environmental
conditions.
rExpand the number of galactic white dwarfs by a factor 10
for which high resolution/high signal-to-noise UV spectra
are available.
rIncrease by a factor 10 the number of binary systems with
white dwarf components for which astrometric masses can
be obtained.
rBe able to study uniform, co-eval populations of white
dwarfs in globular clusters, the Magellanic Clouds and
nearby galaxies.
4. Future white dwarf research in the far ultraviolet
4.1. White dwarfs in the galaxy
A large-scale survey of the hot white dwarfs in the galaxy will
provide observations which can simultaneously address two
broad areas of astrophysics: the local interstellar medium,
its composition, ionization and structure; and the degener-
ate stars, their origin and evolutionary history as well as the
detailed modeling of critical physical processes in their pho-
tospheres. The primary broad scientific objectives would be
following:
rDefine the evolutionary history of the hot white dwarf stars
through detailed modeling of their photospheric composi-
tion and structure.
rStudy the occurrence of circumstellar material surrounding
the white dwarfs and their interaction with the ISM.
rMap out the 3-D structure of the local interstellar medium
(LISM) and determine its composition and ionization state.
rUse the morphology of the LISM and the estimated el-
emental diffusion (gravitational settling) times for white
dwarf photospheres to provide a crucial test of the abil-
ity of interstellar accretion processes to explain abundance
patterns in cooler white dwarfs.
rIdentify and characterise new non-interacting white dwarf
binary systems.
rCarry out a long-term astrometric programme to determine
white dwarf masses.
4.1.1. White dwarf composition
Although there is a qualitative understanding of how
the abundance patterns vary across the hot white dwarf
Springer
Astrophys Space Sci (2006) 303:3–16 13
population, the detailed picture is quite confused and some
very important questions need to be answered. The observed
abundances will reflect the balance of several processes, in-
cluding mass loss, radiative levitation, gravitational settling,
and accretion from the LISM. DA white dwarfs have been
selected for follow-up UV studies mainly on the basis of
their EUV fluxes, low values indicating the presence of pho-
tospheric metals. Thus, the existing observational sample is
strongly biased toward such stars. Few stars having appar-
ently pure-H atmospheres have been observed. In addition,
the existing data are highly non-uniform in wavelength cov-
erage, signal-to-noise and spectral resolution, which yield
non-uniform detection criteria for spectral features. There-
fore, the apparent absence of metals may be as much a func-
tion of the weak detection limits and too few observations of
appropriate stars than a real lack of metals. Hence, we have
no idea whether the small group of heavy element-rich stars
(see Fig. 4) are typical of the cooler group below 50,000 K,
or whether they are truly unusual. Are there really two dis-
tinct groups of stars with and without metals? Or, is this an
artifact of the small sample and in reality compositions range
between the two extremes? It is hard to explain why any DA
would have no heavy elements at all. Thus, establishing the
frequency of pure-H envelopes is an important component of
providing an answer to this problem. In particular, we need
to properly sample the lower temperature white dwarfs, es-
pecially within the 20,000–35,000 K region, which are not
well represented in earlier studies and existing data.
PG1159 stars are rare objects, about 40 are known. Only
a few of them have been studied in detail in the UV. High-
resolution UV observations are essential, because most di-
agnostic metal lines observed in these extremely hot stars
are located in this wavelength region. The wide spread in
element abundances, as well as the observed iron-deficiency
and neon- and fluorine-overabundances show that PG1159
stars have a large, and unique, potential to study mixing and
fusion processes whose consequences are usually unobserv-
able in other stars. As a consequence of a late He-shell flash,
PG1159 stars exhibit intershell matter that normally remains
hidden in the stellar interior. In contrast to DA and DO white
dwarfs, the observed element abundances in PG1159 stars
are not affected by gravitational settling, hence, abundance
patterns still do reflect the history of these stars.
4.1.2. Circumstellar material
What appears to be circumstellar gas has been detected in
most of the white dwarfs observed in high resolution HST
spectra. What is the nature of this material? Is it present in all
white dwarfs with photospheric metals or is there a temper-
ature cut-off? A lower temperature (greater age) limit would
imply that we are looking at a nebular remnant, which dis-
perses with time. Apportioning unresolved lines to circum-
stellar and photospheric components for the whole sample
of stars is essential for correct atmospheric abundance mea-
surements. For example, when observed by HST, the photo-
spheric C abundance of G191-B2B was really found to be a
factor 5 lower than the value obtained by IUE (see Barstow
et al., 2003b).
4.1.3. 3-D structure of the ISM
High-resolution spectra of LISM absorption lines from abun-
dant ionic species (C II, C II, N I, O I, Si II-III, S II, Fe II, Zn
II, Cr II, Mg I-II) will provide several quantitative measure-
ments of nearby interstellar gas. We will be able to measure
the line-of-sight densities and composition of the LISM, de-
rive velocities, to probe the gas kinematics, and determine
the ionization of the gas. It is important that the chosen white
dwarf sample includes enough different lines-of-sight to pro-
vide a true 3-D picture. The high resolution of the echelle data
for these white dwarfs will be instrumental in resolving the
inherent complex velocity structure seen even in the very
local gas within 15 pc (e.g., Lallement et al., 1986; Sahu
et al., 2000a,b), and provide the means to obtain reliable
ionic column densities for the individual velocity compo-
nents in the LISM. Does a single bulk velocity vector fit all
the lines of sight through a particular cloud, or is the gas
fragmented into filamentary structures more characteristic
of low velocity shocks? An adequate sample of white dwarfs
will provide extensive sampling for distances out to 50–100
parsecs.
The problem of the variability of the D/Hratio in the
LISM appears to be on the way to resolution, but there are
many details that still need to be addressed (Moos et al.,
2002; Sahu et al., 2000a,b). Specifically, obtaining reliable
D/Hratios is not easy. Extreme care must be taken in data
reduction and calibration and accounting for multiple veloc-
ity components along the line of sight. One must look at the
heavy ions to determine what velocities are present, because
they can be easily masked in the intrinsically broader profiles
of light HI and DI.
Column densities, measured from high resolution far-UV
spectra, will be useful for determining chemical abundances
in the LISM clouds. Comparing the LISM abundances with
those of Savage and Sembach (1996) would determine if
the cloud has abundances and depletions similar to warm
partially ionized gas observed in the more distant ISM.
Examination of the white dwarf data has revealed an un-
recognized problem, namely the frequent occurrence of in-
terstellar Si III λ1206 in many lines-of-sight. This is diffi-
cult to reconcile, since this ion has a high charge-exchange
rate with neutral hydrogen. Its presence suggests substan-
tial amounts of warm ionized gas have been unaccounted
for in the LISM. It is intriguing that low ionization species
seen in G191-B2B (Sahu et al., 2000a,b) are close to the
Springer
14 Astrophys Space Sci (2006) 303:3–16
velocity of interstellar Si III. An examination of S II/H I
and C II/H I ratios compared to those found for the LIC
at 19.3 km/s, suggests the H ionization fraction is near 0.7
for the 8.6 km/s gas. Another problem is that the short-
ward “circumstellar” components of C IV in G191-B2B
also have a close velocity coincidence with the Si III and
the 8.6 km/s low ionization species. Photoionization calcu-
lations find no way to have all of these ions in the same
gas. High resolution (R50,000–100,000) data for other
sightlines should be able to determine the origin of the
Si III.
4.1.4. Interstellar accretion
For the cooler DAZ stars (6000 <Teff <12,000 K) the ex-
istence of heavy elements such as Ca, Mg and Fe has tradi-
tionally been attributed to interstellar accretion (e.g., Dupuis
et al., 1993). For DA stars at intermediate temperatures
(20,000K<Teff <25,000 K) the presence of heavy ele-
ments poses a dilemma.
In DAZ stars the retention times for high Z elements in
the observable photospheres, while substantial (i.e., 104
yr), remains short compared to the thermal (cooling) time
scale of the photospheres. It is possible to imagine that DAZ
white dwarfs represent stars passing through, or recently
passed through, a diffuse interstellar cloud and having ac-
creted heavy elements. As long as the fraction of DAZ stars
remained small it was possible to entertain such views since
the number of these stars were the result of infrequent en-
counters between diffuse clouds and white dwarfs. However,
Zuckerman et al. (2003) have shown that the occurrence of
cool DAZ stars approaches 25%, which is well in excess of
the fraction that might be attributed to ISM accretion. Like-
wise, other possible explanations involving intrinsically rare
events such as comet impacts are equally untenable. Yet,
Zuckerman et al., were able to demonstrate a correlation be-
tween the presence of low mass companions and the DAZ
phenomena.
Thus, some form of ongoing circumstellar accretion ap-
pears necessary to explain the bulk of the DAZ stars. Typi-
cally the DAZ stars are too faint and too cool to search for the
UV presence of heavy elements. However, it is possible to in-
vestigate the DAZ phenomena in a hotter class of DA stars at
UV wavelengths. In earlier work, there has been little evalua-
tion of the actual conditions of the interstellar medium along
the lines of sight to the known DAZ white dwarfs. Neverthe-
less, knowledge of the morphology of the LISM is necessary
to evaluate and critically test the accretion model, since the
local distribution of interstellar clouds directly determines
the efficiency of interstellar accretion. In general, interstel-
lar features can be distinguished from stellar features on the
basis of velocity and the presence of excited fine-structure
lines. We can use the several density sensitive indicators to
probe the gas phase density of any accreting medium in the
vicinity of the star. For example, an important diagnostic tool
is the presence of the collisionally exited C IIλ1335 line,
from which estimates of the ambient electron density can be
determined (Holberg et al., 1999). The presence of ground
state and exited C I lines indicate the presence of a medium,
which is sufficiently dense to be effectively self-shielding
with respect to UV radiation shortward of the C I ionization
limit. Even if excited lines are not detected, the mere pres-
ence of significant column density, as evidenced by strong
IS absorption, is important.
4.1.5. White dwarfs in binaries
Many new white dwarf binaries with main sequence com-
panions are being discovered by the GALEX UV sky sur-
vey. In these systems the presence of the companion ob-
scures the white dwarf at long wavelengths. Hence, for
these objects, far-UV spectroscopy is essential to deter-
mine the basic physical parameters of the white dwarfs.
For example, temperature and surface gravity can be de-
termined from model atmosphere analyses of the hydrogen
Lyman series lines (e.g., Barstow et al., 2003a), while pho-
tospheric composition can be determined from any heavy
element lines present, as discussed above. Since many
of the brighter companions will be members of the Hip-
parcos catalogue, and therefore have well-determined dis-
tances, the sample of binary white dwarfs can be used to
study the mass-radius relation and the initial-final mass
relation. The latter is extremely uncertain and is impor-
tant in validating potential models for type Ia supernovae
progenitors.
Only about half a dozen of the known binary systems
have sufficiently short orbital periods for astrometric infor-
mation to be obtained on sensible timescale. A subset of
any newly discovered binaries will also fall into this group.
Thus far astrometric orbits and directly determined masses
are only available for three white dwarfs, Sirius B, 40 Eri B
and Procyon B. All other WD masses are based on gravita-
tional redshifts or spectroscopic log gdeterminations, which
require theory-dependent assumptions, or on generally un-
certain measurements of interacting close binaries. Knowl-
edge of the masses is, in turn, vital in testing the theory of
WD structure (e.g., the mass-radius-core composition rela-
tion), understanding the history of star formation in the solar
neighbourhood, and setting limits on the age of the Galaxy.
4.2. White dwarfs outside the galactic disk
Imaging surveys of white dwarfs have been carried out
in globular clusters, but most of the individual stars are
only characterized by broad band photometry. This pro-
vides almost no information on the white dwarf masses and
Springer
Astrophys Space Sci (2006) 303:3–16 15
weak temperature constraints. Knowledge of the white dwarf
masses is essential for determining the cooling age of indi-
vidual stars and interpretation of the observed luminosity
functions. Study of white dwarfs in globular clusters yields a
number of advantages compared to the galactic population.
rAll the stars lie at the same, known distance.
rThe stars are co-eval.
rThe stars are all descendents of a uniform population.
Knowledge of white dwarf temperature and gravity (and,
therefore mass) in globular clusters will provide a direct cali-
bration of the initial-final mass relation. In particular, the up-
per limit of the progenitor mass will be reliably established
for the first time, which has important implications for mod-
els of SNIa systems. In the galaxy most of the white dwarf
progenitors for stars in the disk will probably have had popu-
lation I metallicities. With all progenitors in a globular cluster
being population II, the metallicity and, as result, their prior
evolution will be well determined. Since post-main sequence
evolution is affected by stellar metallicity, in particular in es-
tablishing core He burning, we would expect the resulting
white dwarf population to have different characteristics to
those in the galaxy.
5. The future need for far UV missions
In Section 4 we have outlined in detail the scientific goals
for future studies of white dwarfs. This wealth of science is
only possible through a programme of observations in the far
ultraviolet waveband. The principal need is for high resolu-
tion spectroscopy, but diffraction limited imaging is also of
importance. During the past 15 years, these joint capabilities
have been provided by the Hubble Space Telescope, follow-
ing on from 18 years of operations with IUE. Although of
great importance, the relatively small aperture of IUE limited
high resolution studies of white dwarfs to a handful of the
brightest examples. Using a variety of instruments HST has
provided us with a flow of high signal-to-noise and resolu-
tion (R50,000–100,000) spectra of white dwarfs and the
first diffraction limited imaging of white dwarfs in binary
systems. However, HST time has had to be shared across
a wider range of wavelengths including the visible and IR
bands. Since 1999, this has been complemented by the avail-
ability of the FUSE mission, working down to shorter wave-
lengths than HST (912 ˚
A cf. 1050 ˚
A), but with more modest
spectral resolution (R20,000). Sadly, the STIS instrument
on HST failed in August 2004, ending the UV spectroscopic
capability for the foreseeable future. Also, FUSE is prob-
ably nearing the end of its lifetime. While operations have
been maintained through the heroic efforts of the FUSE team,
continued degradation of its attitude control system (through
gyroscope and reaction wheel failures) will eventually lead
to its termination. Within current mission plans access to far
UV spectroscopy is likely to end soon with no prospect of any
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