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85 is devoted to summarizing our conclusions.
5 is devoted to summarizing our conclusions.
Por most classes of astronomical objects. the distinctive phenomenological characteristics used to identify their members are clearly related to the distinctive physical characteristics that define their fundamental nature.
For most classes of astronomical objects, the distinctive phenomenological characteristics used to identify their members are clearly related to the distinctive physical characteristics that define their fundamental nature.
In the case of SSSs. however. the connection between the phenomenological definition and the gross physical properties is not unique.
In the case of SSSs, however, the connection between the phenomenological definition and the gross physical properties is not unique.
The broad range of temperatures and luminosities that characterize SSSs can be associated with many different types of physical systems.
The broad range of temperatures and luminosities that characterize SSSs can be associated with many different types of physical systems.
Analogous statements apply to the new class of QSSs.
Analogous statements apply to the new class of QSSs.
In this section we provide an overview of physical models for VSSs.
In this section we provide an overview of physical models for VSSs.
SSS-like emission is expected from hot WDs.
SSS-like emission is expected from hot WDs.
Approximately half of the SSSs with optical IDs in. the Milky Way and Magellanic Clouds are systems containing hot white dwarfs: symbioties. recent novae. the central stars of planetary nebulae.
Approximately half of the SSSs with optical IDs in the Milky Way and Magellanic Clouds are systems containing hot white dwarfs: symbiotics, recent novae, the central stars of planetary nebulae.
In most of these systems. the high temperatures and lummosities are associated with episodic nuclear burning. or with cooling subsequent to nuclear burning.
In most of these systems, the high temperatures and luminosities are associated with episodic nuclear burning, or with cooling subsequent to nuclear burning.
The question that was raised by the discovery of binary SSSs like CAL 83 and CAL 87 is whether the soft luminous emission from them and from some of the other mysterious SSSs is due to quasi-steady nuclear burning of matter accreted by a WD from a Roche-lobe-filling companion.
The question that was raised by the discovery of binary SSSs like CAL 83 and CAL 87 is whether the soft luminous emission from them and from some of the other mysterious SSSs is due to quasi-steady nuclear burning of matter accreted by a WD from a Roche-lobe-filling companion.
Conceptually. there is a natural place in the pantheon of accreting WDs. for WDs in contact systems with high-enough accretion rates to permit quasi-steady nuclear burning.
Conceptually, there is a natural place in the pantheon of accreting WDs, for WDs in contact systems with high-enough accretion rates to permit quasi-steady nuclear burning.
Among the accreting WD systems that dominated studies before the discovery of SSSs were cataclysmic variables (CVs) and symbiotics.
Among the accreting WD systems that dominated studies before the discovery of SSSs were cataclysmic variables (CVs) and symbiotics.
In CVs the aceretion rates are low (typically <1077M.. yr!) and the observed luminosities can be explained by accretion power.
In CVs the accretion rates are low (typically $\leq 10^{-9} M_\odot$ $^{-1}$ ) and the observed luminosities can be explained by accretion power.
The accretion rates are low because. although the donor fills its Roche lobe. the mass ratio GMaoo/Mwp) 1s typically small and the donor is typically not evolved.
The accretion rates are low because, although the donor fills its Roche lobe, the mass ratio $M_{donor}/M_{WD}$ ) is typically small and the donor is typically not evolved.
In systems with larger mass ratios. the accretion rate should be larger. in some cases large enough (~1O77M.. yr) that the acereting material would be burned in a quasi-steady manner.
In systems with larger mass ratios, the accretion rate should be larger, in some cases large enough $\sim 10^{-7} M_\odot$ $^{-1}$ ) that the accreting material would be burned in a quasi-steady manner.
Although the appearance of such systems would be quite different from that of normal CVs. especially because the luminosity would be 10—100 times larger than the aceretion luminosity. which would itself be larger than typical. they could represent an epoch of normal CV lifetimes. as follows. (
Although the appearance of such systems would be quite different from that of normal CVs, especially because the luminosity would be $10-100$ times larger than the accretion luminosity, which would itself be larger than typical, they could represent an epoch of normal CV lifetimes, as follows. (
1) When Roche lobe overflow of a main-sequence star more massive than the WD begins. there could be a brief ramp-up time (~10° yrs). during which the mass-transfer rate would be below the value needed for nuclear burning.
1) When Roche lobe overflow of a main-sequence star more massive than the WD begins, there could be a brief ramp-up time $\sim 10^5$ yrs), during which the mass-transfer rate would be below the value needed for nuclear burning.
The system might appear to be a CV. although with a brighter donor star than is typical of CVs. (
The system might appear to be a CV, although with a brighter donor star than is typical of CVs. (
2) The aceretion rate reaches the value needed for quasi-steady nuclear burning. driven by the thermal-time-scale adjustment of the donor to a shrinking Roche lobe.
2) The accretion rate reaches the value needed for quasi-steady nuclear burning, driven by the thermal-time-scale adjustment of the donor to a shrinking Roche lobe.
There ensues an epoch during which high luminosities would be generated by nuclear burning. with &T in the SSS range: the system would be a close- SSS (CBSS).
There ensues an epoch during which high luminosities would be generated by nuclear burning, with $k\, T$ in the SSS range; the system would be a close-binary SSS (CBSS).
The CBSS era ends after ~10 years. roughly the thermal time scale of the donor. when the mass ratio has reversed.
The CBSS era ends after $\sim 10^7$ years, roughly the thermal time scale of the donor, when the mass ratio has reversed.
After the system cools. the SSS behavior ends. (
After the system cools, the SSS behavior ends. (
3) The donor. which is now less massive and which is no longer stressed by a quickly shrinking Roche lobe. may underfill its Roche lobe for some time.
3) The donor, which is now less massive and which is no longer stressed by a quickly shrinking Roche lobe, may underfill its Roche lobe for some time.
Eventually. either due to dissipative processes or to its own evolution and expansion. the donor will again fill its Roche lobe and the system may again appear to be aCV. (
Eventually, either due to dissipative processes or to its own evolution and expansion, the donor will again fill its Roche lobe and the system may again appear to be a CV. (
4) If dissipative processes were responsible for initiating the second epoch of mass transfer. then the system would simply be a normal CV. following a standard evolutionary track.
4) If dissipative processes were responsible for initiating the second epoch of mass transfer, then the system would simply be a normal CV, following a standard evolutionary track.
We also note that (5) If stellar evolution played the key role in re-establishing contact. then the orbital period will increase as mass transfer proceeds.
We also note that (5) If stellar evolution played the key role in re-establishing contact, then the orbital period will increase as mass transfer proceeds.
in some systems. the donor's envelope would be exhausted before the donor became very evolved.
in some systems, the donor's envelope would be exhausted before the donor became very evolved.
In others. the donor would become a full-blown giant.
In others, the donor would become a full-blown giant.
In the latter case. some systems could again experience an epoch of high-mass transfer and SSS-like behavior.
In the latter case, some systems could again experience an epoch of high-mass transfer and SSS-like behavior.
These systems would be wide-binary SSSs (WBSSs). (
These systems would be wide-binary SSSs (WBSSs). (
6) Wide binary SSSs are more commonly expected when the donor first fills its Roche lobe às a giant.
6) Wide binary SSSs are more commonly expected when the donor first fills its Roche lobe as a giant.
Some WBSSs would be virtually indistinguishable from symbiotic systems.
Some WBSSs would be virtually indistinguishable from symbiotic systems.
The donor stars in most symbiotics are thought to feed the WD through winds. rather than through Roche-lobe overflow.
The donor stars in most symbiotics are thought to feed the WD through winds, rather than through Roche-lobe overflow.
Because the donors are very evolved. the rate of wind capture by the WD can be high enough that episodic or quasi-steady nuclear burning can occur.
Because the donors are very evolved, the rate of wind capture by the WD can be high enough that episodic or quasi-steady nuclear burning can occur.
There is good evidence that the high lummosities of symbiotics are powered by nuclear burning. which may be episodic in some systems or steady in others.
There is good evidence that the high luminosities of symbiotics are powered by nuclear burning, which may be episodic in some systems or steady in others.
Thus. quasi-steady-nuclear burning WDs in. which a companion fills its Roche lobe is a natural extension of the class of CVs. with CVs representing the lower-accretion- systems.
Thus, quasi-steady-nuclear burning WDs in which a companion fills its Roche lobe is a natural extension of the class of CVs, with CVs representing the lower-accretion-rate systems.
More massive and/or more evolved donors lead to supersoft binaries.
More massive and/or more evolved donors lead to supersoft binaries.
The Roche-lobe filling SSS binaries
The Roche-lobe filling SSS binaries
The question of what triggers the star formation in the disks of galaxies, and in particular the role of the spiral arms, has been the subject of a long debate.
The question of what triggers the star formation in the disks of galaxies, and in particular the role of the spiral arms, has been the subject of a long debate.
Soon after the theory of spiral arms was proposed by ?),, ?) showed that shock waves are produced in the passage of the gas across the spiral gravitational perturbation, and this could trigger star formation.
Soon after the theory of spiral arms was proposed by\cite*{LinShu69}, , \cite{Roberts69} showed that shock waves are produced in the passage of the gas across the spiral gravitational perturbation, and this could trigger star formation.
This was not a surprising result, since the arms are obvious to us because they are the areas containing bright stars. ?),,
This was not a surprising result, since the arms are obvious to us because they are the areas containing bright stars. \cite{Oort74},
analyzing the gas distribution in M51, suggested that a constant fraction of the gas per passage through the arms is transformed into stars, and consequently that it would be natural to expect the star formation rate to be proportional to Ω—Ωρ.
analyzing the gas distribution in M51, suggested that a constant fraction of the gas per passage through the arms is transformed into stars, and consequently that it would be natural to expect the star formation rate to be proportional to $\Omega-\Omega_p$.
?) tested this expression with a sample of nearby galaxies, by plotting a chemical enrichment indicator ([O 11]/H@) as a function of Q—Ωρ, both quantities being measured at a normalized radius, and concluded that the star formation rate is proportional to Q—Ωρ.
\cite*{Jensen76} tested this expression with a sample of nearby galaxies, by plotting a chemical enrichment indicator ([O $\beta$ ) as a function of $\Omega-\Omega_p$, both quantities being measured at a normalized radius, and concluded that the star formation rate is proportional to $\Omega-\Omega_p$.
However, ?) argued that the correlation found by Jensen et al.
However, \cite{McCall86} argued that the correlation found by Jensen et al.
could be the result of independent processes having similar dependence on galactic radius, and presented observations ofO abundance gradients in two flocculent galaxies, NGC5055 and NGC7793, which were found to be very similar to those of grand design spirals.
could be the result of independent processes having similar dependence on galactic radius, and presented observations of abundance gradients in two flocculent galaxies, NGC5055 and NGC7793, which were found to be very similar to those of grand design spirals.
As the flocculent galaxies were supposed to be free of spiral arms, the arms could not be the cause of the observed gradients.
As the flocculent galaxies were supposed to be free of spiral arms, the arms could not be the cause of the observed gradients.
This argument was later weakened by the fact that a clear spiral structure was found in NGC5055, both in the infrared (?) and in molecular gas (?)..
This argument was later weakened by the fact that a clear spiral structure was found in NGC5055, both in the infrared \citep{Thornley96} and in molecular gas \citep{Kuno97}.
Thornley found infrared spiral structures in several other flocculent galaxies, which is an indication that it is a normal characteristic of these galaxies.
Thornley found infrared spiral structures in several other flocculent galaxies, which is an indication that it is a normal characteristic of these galaxies.
Nevertheless, this does not prove that the spiral arms are responsible for the observed gradients.
Nevertheless, this does not prove that the spiral arms are responsible for the observed gradients.
The nature of the flocculent galaxies is not fully understood, and it is not clear if they can be considered (or not) as good examples of galaxies where stochastic self-propagating star formation dominates.
The nature of the flocculent galaxies is not fully understood, and it is not clear if they can be considered (or not) as good examples of galaxies where stochastic self-propagating star formation dominates.
As a first step, our effort will be to understand the star formation mechanism in normal spiral galaxies.
As a first step, our effort will be to understand the star formation mechanism in normal spiral galaxies.
?) investigated theO abundance profile of seven nearby galaxies and found a distinct change in slope for 3 of them.
\cite*{Zaritsky92} investigated the abundance profile of seven nearby galaxies and found a distinct change in slope for 3 of them.
These changes cannot be associated with calibration problems, since precisely the same methods were used in all the cases.
These changes cannot be associated with calibration problems, since precisely the same methods were used in all the cases.
Later ?) extended the study to a larger number of galaxies, and found for a few of them not only a change of slope, but even possible reversals, with metallicity increasing at large radii (see for instance the profiles of NGC3369).
Later \cite*{Zaritsky94} extended the study to a larger number of galaxies, and found for a few of them not only a change of slope, but even possible reversals, with metallicity increasing at large radii (see for instance the profiles of NGC3369).
It should be noted that a reversal is easily explained if the star formation rate is considered to be proportional to |Q— Q,|, whichis the same expression above but takes
It should be noted that a reversal is easily explained if the star formation rate is considered to be proportional to $|\Omega-\Omega_p|$ , whichis the same expression above but takes
Bl has a proper motion of 0.21 mas yr!. which is roughly similar to previous observations.
B1 has a proper motion of 0.21 mas $\rm yr^{-1}$, which is roughly similar to previous observations.
In the meantime. component B3 also experienes a linear expansion with an apparent velocity of 0.32 mas vr.
In the meantime, component B3 also experienes a linear expansion with an apparent velocity of 0.32 mas $\rm yr^{-1}$.
Although this model can explain the motion of the components. there is another possibility.
Although this model can explain the motion of the components, there is another possibility.
In the second case. component B2 was accelerated after remaining. almost stationary at à core-separation of about 0.5 mas for several years. while component B3 decelerates to a significantly low velocity at the same core separation of about 0.5 mas.
In the second case, component B2 was accelerated after remaining almost stationary at a core-separation of about 0.5 mas for several years, while component B3 decelerates to a significantly low velocity at the same core separation of about 0.5 mas.
Then we applied a helical model under the assumption of three conservative quantities (the jets kinematic energy. angular momentum. and the momentum along the jet axis: Steffenetal. 1995) to interpret the kinematics of the components.
Then we applied a helical model under the assumption of three conservative quantities (the jet's kinematic energy, angular momentum, and the momentum along the jet axis; \citealp{Stef095} 1995) to interpret the kinematics of the components.
The solutions for each component can be used to explain the projected trajectories and the apparent velocity with time as well as to demonstrate in particular the low apparent velocity of the two components B2 and B3.
The solutions for each component can be used to explain the projected trajectories and the apparent velocity with time as well as to demonstrate in particular the low apparent velocity of the two components B2 and B3.
Based on the predictions from the helical model. the components may have the same mode of motion.
Based on the predictions from the helical model, the components may have the same mode of motion.
Future observations are needed to provide new kinematical constraints for the motion of these three components in this source.
Future observations are needed to provide new kinematical constraints for the motion of these three components in this source.
the orbital inclination axis. which is furthermore co-aligned with the stellar rotational axis.
the orbital inclination axis, which is furthermore co-aligned with the stellar rotational axis.
We justify these assumptions by the following argument: Close-in planets such as hot Jupiter seem to be tidally locked with their parent stars (c.f.
We justify these assumptions by the following argument: Close-in planets such as hot Jupiter seem to be tidally locked with their parent stars (c.f.
Shkolnik et al.
Shkolnik et al.
2005; Knutson et al.
2005; Knutson et al.
2007).
2007).
À tidal lock hypothesis suggest that the orbital inclination is co-aligned with the planetary rotation axis and with the stellar rotation axis.
A tidal lock hypothesis suggest that the orbital inclination is co-aligned with the planetary rotation axis and with the stellar rotation axis.
The latter assumption was confirmed by the aforementioned spin-orbit measurements of transiting planets.
The latter assumption was confirmed by the aforementioned spin-orbit measurements of transiting planets.
The two contributions vy) (Equ. 9))
The two contributions $v_{\rm \star,p}$ (Equ. \ref{E2:rot}) )
and vy, (Equ. 10))
and $v_{\rm proj, \rm p}$ (Equ. \ref{E2:rot111}) )
to the broadening are finally summed up to Table 1. summarizes the parameters of the planet and its host star.
to the broadening are finally summed up to Table \ref{tab:tauboo} summarizes the parameters of the planet and its host star.
We note in passing that the system also contains a faint M-dwarf component at à separatation of =230AU from the primary (Patience et al.
We note in passing that the system also contains a faint M-dwarf component at a separatation of $\approx 230~\rm{AU}$ from the primary (Patience et al.
2002: Eggenberger etal.
2002; Eggenberger etal.
2003).
2003).
Knowing the stellar mass M,. the planetary minimum mass Mysini. and the RV semi-amplitude of the reflex motion of the star Ky. the RV semr-amplitudes of the planet can be constrained to ranging from A,=0 to 155.6418.3kmsl (Equation 6:; Figure 1)).
Knowing the stellar mass $M_{\star}$, the planetary minimum mass $M_{\rm{p}}\sin i$, and the RV semi-amplitude of the reflex motion of the star $K_{\star}$, the RV semi-amplitudes of the planet can be constrained to ranging from $K_{\rm{p}} = 0$ to $155.6\pm18.3~\rm{km~s^{-1}}$ (Equation \ref{equ:doppler}; Figure \ref{F6:rv}) ).
Due to the absence of transits in high-precision photometry (Henry et al.
Due to the absence of transits in high-precision photometry (Henry et al.
2000). we can constrain the range of possible orbital inclinations.
2000), we can constrain the range of possible orbital inclinations.
The mmimum orbital inclination { that a transit event occurs can be calculated by where Κι and A, is the stellar and the planetary radius. respectively. and & the semi-major axis.
The minimum orbital inclination $i$ that a transit event occurs can be calculated by where $R_\star$ and $R_{\rm p}$ is the stellar and the planetary radius, respectively, and $a$ the semi-major axis.
Using the parameters listed in Table |. and the assumption that Ay=1.2 Ry. we can exclude orbital inclinations of/>82° for τ Boo b. Baliunas et al. (
Using the parameters listed in Table 1, and the assumption that $R_{\rm p}=1.2~{\rm R_{\rm Jup}}$ , we can exclude orbital inclinations of $i>82^{\circ}$ for $\tau$ Boo b. Baliunas et al. (
1997) found that the star 7 Boo rotates rapidly with a period commensurate with the orbital period of the planet. suggesting tidal locking.
1997) found that the star $\tau$ Boo rotates rapidly with a period commensurate with the orbital period of the planet, suggesting tidal locking.
This hypothesis seems to be very likely for close-in planets (e.g. Shkolnik et al.
This hypothesis seems to be very likely for close-in planets (e.g. Shkolnik et al.
2005: Knutson et al.
2005; Knutson et al.
2007). but has not been observationally confirmed so far.
2007), but has not been observationally confirmed so far.
A tidal lock enables us to place an estimate on the orbital inclination (Equation 13)) under the assumption that the stellar equator and the orbital plane are co-aligned.
A tidal lock enables us to place an estimate on the orbital inclination (Equation \ref{E6:rot1}) ) under the assumption that the stellar equator and the orbital plane are co-aligned.
In order to caleulate the orbital inclination. we need to solve where v, is the measured rotational velocity of 7 Boo Qweps=v.sini149 kms!) for the unknown orbital inclination 7. and v the true. but unknown rotational velocity.
In order to calculate the orbital inclination, we need to solve where $v_\star$ is the measured rotational velocity of $\tau$ Boo $v_{\rm proj,\star} = v_\star \sin i = 14.9~{\rm km~s^{-1}}$ ) for the unknown orbital inclination $i$, and $v$ the true, but unknown rotational velocity.
This rotational velocity v, can be estimated by way of Equation 8.. which requires the knowledge of the stellar radius and stellar rotation period.
This rotational velocity $v_\star$ can be estimated by way of Equation \ref{E2:rot1}, which requires the knowledge of the stellar radius and stellar rotation period.
Adopting the parameters listed in Table |.. we determine the maximum rotational velocity of the star to be vy,=20.0+10kms!.
Adopting the parameters listed in Table \ref{tab:tauboo}, we determine the maximum rotational velocity of the star to be $v_\star=20.0\pm 10~{\rm km~s^{-1}}$.
This gives an orbital inclination of#=467 ?,
This gives an orbital inclination of $i=46^{+36}_{-16}~^{\circ}$ .
The resulting mass of the companion to r Boo of Mp=5.77%My wouldclearly confirm the planet hypothesis.
The resulting mass of the companion to $\tau$ Boo of ${M_{\rm p}=5.7^{+2.6}_{-1.5}~~M_{\rm Jup}}$ wouldclearly confirm the planet hypothesis.
We note that the error values of/ and M, correspond to the maximum error range.
We note that the error values of $i$ and $M_{\rm p}$ correspond to the maximum error range.
cm.
cm.
Each plate is viewed by a PAL tube from above and [rom below.
Each plate is viewed by a PM tube from above and from below.
Four identical modules. Lia? in total area and arrangecl side by side in the same horizontal plane. were built altogether.
Four identical modules, $1\,{\rm m^2}$ in total area and arranged side by side in the same horizontal plane, were built altogether.
We readily see that the system. is asymmetric with respect to the vertical direction. thus permitting one to judge of the daemon fux direction and to draw conclusions on some properties of the radiations measured.
We readily see that the system is asymmetric with respect to the vertical direction, thus permitting one to judge of the daemon flux direction and to draw conclusions on some properties of the radiations measured.
In the absence of a αν of particles causing scintillations. there should. be no correlation between the signals of the PM tubes viewing the top and the bottom plates.
In the absence of a flux of particles causing scintillations, there should be no correlation between the signals of the PM tubes viewing the top and the bottom plates.
In these conditions. the distribution IN»GN) of the fairly rare (noise etc.)
In these conditions, the distribution $N_2(\Delta t)$ of the fairly rare (noise etc.)
signals of the second. (bottom) PM tube in the time shift Af between their beginning and. the beginning of the signal from the first (top) PM tube can be approximated with a constant.
signals of the second (bottom) PM tube in the time shift $\Delta t$ between their beginning and the beginning of the signal from the first (top) PM tube can be approximated with a constant.
An exception to this are the cosmic rays producing a strong maximum at Af=0. with the signals appearing [requently simultaneously. in several modules.
An exception to this are the cosmic rays producing a strong maximum at $\Delta t = 0$, with the signals appearing frequently simultaneously in several modules.