T Tauri Protostar

TTauri

T Tauri stars are thought to be young protostars. They are of spectral class G (or more generally F to M) and
have about the same mass of the Sun - thus they are protostars that eventually become yellow dwarf main
sequence stars. The Sun may well have begun life as a T Tauri star.

Stars form from cool clouds of gas (mostly hydrogen and helium) in space, called
nebulae, that are say about
a thousand times as massive as the Sun. If these clouds are cool enough and dense enough, then they
contract under their own gravity and fragment. Each fragment may then continue to collapse into a protostar.
Initially, the clouds collapse
isothermally - meaning that their temperature remains more or less the same
(cool) because although they convert gravitational energy into radiation, this radiation does not heat the cloud
but instead dissociates hydrogen molecules into hydrogen atoms and then ionises the hydrogen and helium
atoms.

Gravitational energy: when you hold an apple in the air, it has gravitational potential energy stored in it. In
lifting the apple up,  you have done work against the Earth's gravitational field which pulls on the apple, giving it
weight. When you release the apple it falls to the ground, faster and faster until it hits the ground. This energy
of movement, or kinetic energy, came from its stored gravitational potential energy. Potential energy is stored
energy that has the potential to do work. Gravitational energy is energy resulting from a gravitational field, so
gravitational potential energy is simply stored gravitational energy. When a huge gas cloud in space contracts
under its own gravity, it is releasing its stored gravitational potential energy.

Ion: an ion is an atom that has lost or gained one or more electrons to gain an overall electric charge. Click this
link to learn more about
atoms and electrons.

Once the gas is ionised, it starts to heat up as it continues to contract (the gravitational potential energy now
heats the gas as the ionised gas is opaque and the energy cannot escape but heats the gas). Eventually a hot
protostar is formed. Once the temperature in the hot core of the star exceeds about one million kelvins,
nuclear fusion begins and a star is born. The vast amounts of energy released by nuclear fusion
generates a pressure that pushes outwards. The pressure increases until it equals the inward-acting force due
to gravity - the star is in a stable balance or equilibrium (called
hydrostatic equilibrium). Stars are actually
self-regulating, and they pulse in and out slightly as they adjust the rate of their nuclear reactions to maintain
this balance. The star will also rotate, since the cloud of gas generally has a slight rotation to it, but just as a
spinning ice skater spins faster by pulling in their arms, so the gas spins faster as it contracts.

As the material falls inward onto the developing protostar, it tends to form an
accretion disc around the star's
equator. This is a disc of material that slowly spirals onto the star, heating up as it does so. Some material also
gets blasted away, however, if it has too much angular momentum (that is if it rotates too fast it gets thrown off).
This material forms
jets, one coming from each pole of the star, and also some material is lost from the edge of
the disc, possibly as a (more or less) spiral wind (called a centrifugal wind) as shown in the diagram above. As
the protostar heats up it also generates a radiation pressure that blasts much of the material away in the form
of a very strong stellar wind, called a
T Tauri wind. Thus, some material comes together to form a star, but
some is also blown away. Indeed, protostars are shrouded in gas, but when the star switches on its nuclear
fusion, it quickly blasts away this shroud and becomes a free star, like the Sun. The jets cease and what
remains of the disc may fragment into rings and eventually form planets.

Young stars, like T Tauri stars, usually rotate rapidly (as the ice skater has just pulled their arms in a very long
way!) and this is probably generates a very strong magnetic field in these stars. This field may interact with the
disc (by dragging against it) and eventually put the breaks on and stop the star spinning so fast that it flies
apart!

See also
flare stars for a description of young stars of lower mass (red dwarfs of spectral class M or K) and Of
stars for massive young hot blue stars.

The definition of T Tauri stars is often taken to include all young dwarf stars of spectral class F, G, K and M and
less than twice the Sun's mass (2 solar masses). Similar young stellar objects (YSOs) of spectral class A and B
are called HAeBe stars (Herbig Ae/Be stars where e stands for emission spectrum, indicating that these stars
are surrounded by clouds of hot and luminous gas which are presumed remnants of the nebula from which the
new stars were assembled). HAeBe stars have a mass of 2 to 8 times the Sun's mass (2 to 8 solar masses).
Class O stars probably pass through similar youthful stages, perhaps Of stars.

All these YSOs are between about 1 to 10 million years old. YSOs are thought to pass through four
evolutionary stages corresponding to objects of class 0 to III:

    
Class O objects - obscured by a spherical, infalling dust cloud (a molecular cloud of 99% gas and 1%
dust), forming a cool envelope of 15-30 K and emitting radio waves. The object is not yet a star or YSO.
    
    
Class I objects - these young stars are undergoing gas accretion (gas is falling onto their surface as they
grow in mass). An accretion-driven wind drives away envelope gas from the poles and a YSO is born. Most of
the star's light is absorbed and re-radiated by the envelope at long wavelengths. This stage lasts about 100
000 years. A circumstellar accretion disc (protoplanetary disc) is forming.

    
Class II object - the YSO is surrounded by an accretion disc with a similar composition to the original
molecular cloud (99% gas, 1% dust) and possibly optically thick to begin with (opaque) due to the dust
scattering light. This matter heats up and emits mostly in the infrared and optical. An accretion disc forms
around the equator of the forming star, in place of the original accreting spherical cloud (stage 0), due to
conservation of angular momentum - material far from the star has large angular momentum and must lose this
momentum as it approaches the star, spinning faster as it gets closer (like the spinning ice-skater who pulls in
her arms). Matter falling in towards the poles end-up spinning too fast and would need to lose almost all of its
angular momentum, so instead is spun-away in polar jets.

    
Class III object - The disc drains as its matter accretes onto the central star and the star's radiation bathes
the thinning disc, evaporating away dust grains as the disc heats up and becomes optically thin. Gas
diminishes as the disc drains and accretion is much reduced if not absent and most of the spectrum is
dominated by light from the star's photosphere, which now shines clear of the dispersing nebula. Surrounding
dust and gas still contributes some infrared light to the spectrum. The YSO is now moving on to the main
sequence as a fully-fledged star.

T Tauri stars are YSOs of low-mass stars of class I and II. They have starspots and powerful stellar winds and
emit strongly in the X-ray and radio wavelengths and are variable stars (i.e. they show dramatic fluctuations in
their luminosity). HAeBe stars are also class I and class II objects, but of intermediate to massive A and B class
stars.

Protoplanetary Discs

Outside the accretion disc, in both T Tauri stars and HAeBe stars, is an extended dusty torus of matter which
feeds the accretion disc in the initial stages. Dust grains form in this cooler torus, crystallising and accumulating
matter, and when the dust grains become large enough they are no longer swept along with the gas, but settle
out to form a dark, dusty disc inside the torus. The torus becomes depleted of gas, both by feeding gas into the
accretion disc, and by photoevaporation (the intense light of the growing star blasts away the gas in an intense
stellar wind). These dust grains may eventually coagulate into planets, hence the disc around a YSO is also
called a
protoplanetary disc. The structure of a typical protoplanetary disc is shown below:

Protoplanetary disc

The colours of these discs varies according to the size of the dust grains. Smaller grains scatter light by
Rayleigh scattering (as in the Earth's atmosphere) and colour the material blue, whilst larger grains scatter
redder light. Since regions nearer to the forming star are hotter, they may contain smaller particles as the
dust grains evaporate (though smaller grains evaporate first). Hence we have depicted the
inner
accretion disc
in blue, the dusty torus in orange-red. The accretion disc may be optically thick or thin, and
is probably more opaque in the earlier stages, but may be a pure gas disk. The inner disc is hot, at a
temperature much greater than 1000 K and about 0.1 to 1 AU in radius. Some dust may form in regions of
the disc shadowed by optically thick gas (dust plume).

AU:
astronomical unit, the average distance between the Sun and Earth, 1 AU = 149 597 870 km ~ 150
million km.

At some distance from the star, depending on the star's age and size, temperatures drop below about 1500
K and the disc gives way to a
thick extended torus or thick outer dusty disc. At 1500K dust
sublimates/evaporates, so when temperatures drop below this point dust suddenly forms and observations
suggest that at this
dust sublimation temperature the forming dust forms a thick, but smoothly rounded
wall or ring around the star (the dusty inner rim). Behind this thick rim extends a disc of gas and dust, flaring
out away from the star, being sculpted by the stars winds blasting over the top of the inner rim. The torus
may exhibit a greater or lesser degree of
self-shadowing, especially when it is contains sufficient material
and dust. In some systems the torus may be shallower, flatter and less flared.

Flared discs may be classified as Group 1 discs, whilst flat discs, which are more self-shadowed, are Group
II discs.  These different geometries are predicted from different infrared (IR) emissions ( having a flat,
double-peak between 2 to 100 micrometres wavelength in Group I, weaker declining IR emission from 2 to
100 micrometres in Group II).

Close to the star, the infalling matter in the accretion disc may move out from the disc, along magnetic field
lines which guide the accreting material onto the star's surface, largely toward the polar regions. This
magnetospheric accretion, driven by magnetic forces, is more likely to happen in stars with stronger
magnetic fields and when rates of accretion are low. Accretion rates may well vary periodically, as in any
accretion disc, undergoing periods of stable and low rate accretion, which is likely to be magnetically driven,
and brief episodes of rapid accretion, in which the ram pressure of infalling matter might be strong enough
to overcome the magnetic field, in which case material from the disc will fall directly onto the star's equator,
causing the star to dramatically brighten for a time.
Shock waves may emit intense radiation (UV and
X-ray) when the matter hits the star's surface. The region of magnetosphere-dominated accretion is
expected to be about 0.03 AU in radius and gives rise to a continuum (i.e. extending over a range of
wavelengths) of UV emission.

The torus can resupply the accretion disc. However, at some point the two become disconnected, possibly
by the formation of large planets which create gaps in the outer dusty disc, gathering up any matter that
attempts to cross their orbit. The supply of gas may also run out as gas is accreting, as gas is blasted away
by the intense stellar winds of the new star, and by coagulating into dust grains, planetoids and planets.
The presence of large planets of one or more Jupiter masses, seems to be in part responsible for
distortions of the discs, including warping of the disc surface, departures from circularity as the disc become
elliptic, and displacement of the disc, so that the star is not at the exact centre. When accretion slows and
the accretion disc becomes optically thin, the radiation from the star is free to pass through it and to blast
the dusty outer disk directly, and the inner region of this disc evaporates away as a widening hole forms at
the centre of the protoplanetary disc. The outer parts of the disc may remain to condense into planetoids,
as in the Kuiper belt of the Solar System. The outer disc is a
reservoir of gas and dust and cold at a
temperature of about 10 to 30 K.

The stages of disc development can be summarised as follows:

    Stage 1 - a massive flared outer disc and inner accretion disc, possibly with magnetospheric accretion.
    Stage 2 - Ultraviolet (UV) and X-rays from the central star evaporate and blast away the upper and lower
surfaces of the disc.
    Stage 3 - A settled disc, in which dust grains of increasing mass settle out to form a thinner dust ring in
the mid-plane of the disc.
    Stage 4 - Accretion has stopped and the disc photoevaporates from the inside-out.
    Stage 5 - A
debris disc remains, depleted of gas, formed by further coagulation of dust grains in the
disc mid-plane. This disc may contain asteroids, planetoids and planets.

The whole cycle from birth to planet formation is expected to take 1 - 1o million years, and planets must
form before matter is depleted. The planet-forming region is thought to extend from the inner dusty rim to
about 10 to 30 AU.

An example of a T Tauri star is TW Hya (TW Hydrae), which has a disc that is elliptical and warped between
about 70 and 140 AU, presumably due to the formation of large planets in this region, and circular further
out. The central star is a spectral class K8 variable star (an orange dwarf).

Planet Formation

The coagulation of gas to form dust grains is predicted to be quite easy. Once the grains reach a threshold size they are also predicted to rapidly enlarge, though this is countered by collisions between grains and fragmentation, which may generate gas even in old discs. A particular problem occurs in achieving sizes greater than 1 meter and models may be missing certain parameters that make this transition easier. Maybe planets only form in a minority of debris discs, on the other hand planet formation may be almost universal.

Initially, many planets may form, but their orbits will be unstable to begin with and some will collide with other planets and some may become ejected from the system, and yet others may spiral into the central star to be consumed. The question remains how many are likely to survive in stable orbits? However, not only do many YSOs have visible protoplanetary discs, but many mature systems are now known to have planets, so despite these difficulties, planets are probably quite common.

Planets are expected to form at distances of about 0.1 to 30 AU from the central star. Planets are also expected to undergo periods of inward migration and some mechanism must break them if they are not to be destroyed by their star. This inward migration might account for the number of 'hot Jupiters' detected - large gas giants orbiting very close to their star.

Binary Systems

In close binary systems, containing two stars orbiting their common centre of mass at separations of about 5 to 100 AU, observations suggest that each star may have its own protoplanetary disc, and a common ring encloses both discs further out. In these systems, we might expect planets to form in orbits around each star and also possibly further out in a single orbit around both stars. The orbital dynamics become complicated and many planets may become ejected from such systems, but some may persists in stable orbits.

O Stars

Planetary systems around O stars seem to be much shorter lived, perhaps less than 1 million years, and so are harder to detect. In many cases it is a possibility that they never form. Massive stars are short-lived and
many will become
neutron stars / pulsars. The mass ejected in the preceding supernova explosion may coagulate and undergo a new phase of planet-building, and some old planets may survive if they were far enough out initially (though they may be ejected or destroyed). Planets are certainly known to exist around neutron stars. O stars stand a good chance of becoming stellar black holes and planets may be able to form around these structures too.

There has been concern that the presence of O stars, with their intense UV radiation could prevent disc formation around nearby stars by photoevaporating their discs too rapidly. However, calculations suggest that only the outer edges of discs in nearby star systems will be eroded, unless the young star is closer than 0.1 pc or about one-third of a light-year (i.e. relatively very close) to the O star. Thus O stars are unlikely to perturb planet formation except in their own systems.

Protoplanetary disc

Above: a model of a protoplanetary disc seen face-on, showing the flat inner blue accretion disc and redder outer disc or dusty torus. Distortions or asymmetries in the outer disc suggest that planets may be forming inside.

Protoplanetary disc

Above: the disc seen at an inclination of 15 degrees.


Example: TW Hydrae - protoplanetary disc

TW hydrae is a T-Tauri star in the constellation Hydra, 176 light years from Earth. The protoplanetary disc of gas and dust orbiting this star is imaged below (NICMOS - Near Infrared camera and Multi-Object Spectrometer). This image has been color-coded by brightness in shades of red (original in b&w). The central dark area is an artifact of imaging and not an actual clearing in the disk (it is due to the optical systems coronagraph which blocks out the very bright light from the central star to enable better imaging of the fainter disc).

Protoplanetary disc
image Credits: NASA , ESA , J. Debes (STScI ), H. Jang-Condell (University of Wyoming), A. Weinberger (Carnegie Institution of Washington), A. Roberge (Goddard Space Flight Center), G. Schneider (University of Arizona/Steward Observatory), and A. Feild (STScI /AURA ) 2005.

Note the gap in the protoplanetary disc which is 1.9 billion miles (3 billion km) wide and is thought to be due to the orbit of an unseen planet sweeping up material as it gradually clears its orbit.

Protoplanetary disc
Image credits: NASA , ESA , and J. Debes (STScI) 2015/16

Above:images taken one year apart reveal a shadow moving counterclockwise across the disc. The shadow takes about 16 years to rotate around the disc and in this image spans about 20o. This shadow is due to a warp in the disc which is thought to be caused by the gravitation pull of an unseen planet orbiting the central star and pulling up material to create a bulge in the disc as it goes.

TW Hydrae is X-ray bright and measurements of plasma temperatures can be explained by a model in which most of the X-rays are due to accretion of mass from the disc onto the central star at a rate of about 10-8 solar masses per year (Kastner and Weintraub, 2002). This star is about 10 million years old and is already showing signs of evolution both with the apparent existence of planets (presumably newly formed) and also with a probable reduction in accretion rate compared to more rapidly accreting younger T Tauri stars of 1 million years of age (Muzerolle et al. 2000).

Example: XZ Tauri - a young binary system

XZ Tauri is a young binary star (about 1 My old) and is shown here emitting a hot bubble of plasma from its northern pole. Presumably accretion from unseen accretion disc(s) around both stars or separately around each is driving this ejection of matter from the system's poles (matter that could not lose sufficient angular momentum to accrete resulting in polar outflow). This binary system consists of an M3 variable T Tauri star in the south (XZ Tau A) and a cooler companion in the north (Krist et al. 1999) - XZ Tau B. It is located in the Taurus constellation and is 140 pc from Earth.

Protoplanetary disc
Image credits: NASA , John Krist (Space Telescope Science Institute ), Karl Stapelfeldt (Jet Propulsion Laboratory), Jeff Hester (Arizona State University), Chris Burrows (European Space Agency /Space Telescope Science Institute )

This outflow could be a discrete (but probably intermittent) eruption or a continuous jet and is thought to originate from one of the stars which emits bipolar jets (Krist et al. 1999). As it travels away from the star, the bubble is expanding sideways at about 33 km/s (corresponding to an internal temperature of about 80 000 K - sideways expansion is due to pressure within the bubble which travels at the speed of sound within the bubble which allows its temperature to be estimated). In 1995 the edge of the bubble was a similar brightness to the interior, suggesting they were at the same temperature. In 1998, however, note that the edge of the bubble has brightened. This is thought to suggest cooling of the outer layers to a temperature cool enough to allow electrons and ions to recombine within the plasma, emitting light (as the electrons lose energy) (Krist et al. 1999).

More recently a larger but fainter southern counterbubble has also been detected and a more recent analysis (Krist et al. 2008) of compact emission knots within each broad bubble suggests that both stars in the binary are emitting bipolar jets. The expanding material in these jets merges together to form the bubble. The northern outflow appears to be a succession of bubbles.  Thus, it would appear that both stars are emitting bipolar jets of varying intensity, resulting in the appearance of a succession of bubbles erupting from each pole.

The faster and dominant jet is generally that from the main star, XZ Tau A at the southern pole. The proposed model to explain the bubble formation is pulsing of the jet (Krist et al. 2008). It is estimated that around 1980 (observer time) a large velocity pulse in the northern jet of XZ Tau A overtook slower ejecta to generate a shock front that formed the compact fireball we can see.

Example: Herbig-Haro jet HH 24

Herbig-Harrow (HH) objects are small bright nebulae of concentrated gas and dust lit-up by radiation or mass outflow from T Tauri stars. They are frequently aligned along the axis of a bipolar outflow to form a Herbig-Haro jet as in HH 24 shown below.

Herbig-Haro jet
NASA and ESA ; Acknowledgment: NASA , ESA , the Hubble Heritage (STScI /AURA )/Hubble-Europe (ESA ) Collaboration, D. Padgett (GSFC), T. Megeath (University of Toledo), and B. Reipurth (University of Hawaii).

Whilst the fireball seen in XZ Tauri is a HH object, what we have with HH 24 is a more classic HH jet. One possible mechanism for the formation of bipolar jets in these systems is funneling of the plasma in the intense stellar winds of these newborn stars by magnetic fields in the accretion disc. If the magnetic field of the disc is more-or-less parallel to the star's rotation axis then the stellar winds will be collimated into two polar jets (Kwan and Tademaru, 1988).


References

Kastner, J.H. and Weintraub, D.A. 2002. Evidence for Accretion: High-Resolution X-Ray Spectroscopy of the Classical T Tauri Star TW Hydrae. The Astrophysical J. 567: 434-440.

Krist, J.E., Stapelfeldt, K.R., Burrows, C.J., Hester, J.J., Watson, A.M., Ballester, G.E., Clarke, J.T., Crisp, D., Evans, R.W., Gallagher III, J.S., Griffiths, R.E., Hoessel, J.G., Holtzman, J.A., Mould, J.R., Scowen, P.A. and Trauger, J.T. 1999. Hubble Space Telescope WFPC2 imaging of XZ Tauri: time evolution of a Herbgig-Haro bow shock. The Astrophysical J. 515: L35-L38.

Krist, J.E., Stapelfeldt, K.R., Hester, J.J., Healy, K., Dwyer, S.J., and Gardner, C.L. 2008. A multi-epoch HST study of the Herbig-Haro flow from XZ Tauri. The Astrophysical J. 136: 1980-1994.

Kwan, J. and Tademaru, E. 1988. Jets from T tauri stars: spectroscopic evidence and collimation mechanism. The Astrophysical J. 332:L41-L44.

Muzerolle, J., Calvet, N., Briceno, C., Hartmann, L. and Hillenbrand, L. 2000. Disk accretion in the 10 Myr old T Tauri stars TW Hydrae and Hen 3-600A. The Astophysical J. 535: L47-L50.


Article updated: 3rd April 2020

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