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Thereore. studying the genesis of local SSCs can vield information about the environment in which globular clusters formed in the early universe. | Therefore, studying the genesis of local SSCs can yield information about the environment in which globular clusters formed in the early universe. |
The conditions required. [or massive star cluster formation are far from understood. but il seems clear that extreme environmeils. uncommon in the local universe. are necessary. | The conditions required for massive star cluster formation are far from understood, but it seems clear that extreme environments, uncommon in the local universe, are necessary. |
Current theories sugeest Chat hieh pressures. such as those due to Luge-scale shocks in merging galaxies (Ehneereen&Elremov1997).. are required in order (o form bound massive star clusters. | Current theories suggest that high pressures, such as those due to large-scale shocks in merging galaxies \citep{elmegreen97}, are required in order to form bound massive star clusters. |
In fact. there is a large body of evidence that globular clusters are formed in galactic interactions | In fact, there is a large body of evidence that globular clusters are formed in galactic interactions |
axisvmnmetirv apply al most only (o mean flow quantities. averaged over sufficient time or space (e.g. Balbus οἱ al. | axisymmetry apply at most only to mean flow quantities, averaged over sufficient time or space (e.g. Balbus et al. |
1994). | 1994). |
The local violation of svuumetries for a turbulent disc. leads lo an intrinsic variability (Blackman 1993). | The local violation of symmetries for a turbulent disc, leads to an intrinsic variability (Blackman 1998). |
We suggest that an additional consequence of the fluctuations may be a local "rocket effect.” | We suggest that an additional consequence of the fluctuations may be a local “rocket effect.” |
When a wind is expelled from the surface of a cise it exerts a force on that surface which compensates for the momentum of the outflow. | When a wind is expelled from the surface of a disc it exerts a force on that surface which compensates for the momentum of the outflow. |
In idealized models of symmetric accretion discs the forces exerted on the dise by the winds on either side of the disc cancel out. | In idealized models of symmetric accretion discs the forces exerted on the disc by the winds on either side of the disc cancel out. |
In reality. although these forces must cancel globally if the dise center of mass is stationary. turbulence induced asvinmetry makes it unlikely that these forces cancel out locally. | In reality, although these forces must cancel globally if the disc center of mass is stationary, turbulence induced asymmetry makes it unlikely that these forces cancel out locally. |
Thus. different portions of the disc can be displaced above or below the initial mean svmnmeltry plane. | Thus, different portions of the disc can be displaced above or below the initial mean symmetry plane. |
Here we investigale how disc turbulence can cause a local imbalance between the winds emanating from (he top aud bottom by averaging the local effect over azimuth. | Here we investigate how disc turbulence can cause a local imbalance between the winds emanating from the top and bottom by averaging the local effect over azimuth. |
The net elfect is to induce a stochastic wander of the vector perpendicular to the plane determined by the azimuthal average of the asvimmeltry for each eddy al a given radius. | The net effect is to induce a stochastic wander of the vector perpendicular to the plane determined by the azimuthal average of the asymmetry for each eddy at a given radius. |
We model this as Brownian motion. aud calculate the associated diffusion coefficient. | We model this as Brownian motion, and calculate the associated diffusion coefficient. |
In Sec. | In Sec. |
2 we discuss the needed basics of Brownian motion. | 2 we discuss the needed basics of Brownian motion. |
li Sec. | In Sec. |
3 we then derive the equation for the aforementioned stochastic diffusion coefficient for (ilt as a function of disc parameters. | 3 we then derive the equation for the aforementioned stochastic diffusion coefficient for tilt as a function of disc parameters. |
In Sec. | In Sec. |
4 we apply this diffusion coefficient to the wobbling of a thin accretion dise and a stochastic wander of the associated jel axis. assumed {ο be parallel to the discs primary angular velocity vector. | 4 we apply this diffusion coefficient to the wobbling of a thin accretion disc and a stochastic wander of the associated jet axis, assumed to be parallel to the disc's primary angular velocity vector. |
We apply the formalism to four astrophysical eneine classes (blazars. voung stellar objects (YSOs). planetary nebulae. microquasars) ancl estimate the wander angles for which the wobble could be observable in these svstems. | We apply the formalism to four astrophysical engine classes (blazars, young stellar objects (YSOs), planetary nebulae, microquasars) and estimate the wander angles for which the wobble could be observable in these systems. |
As a particular example. we compare the crudely predicted wobble to that observed in $5433 and fined agreement (ο order of magnitude for a plausible choice of disc parameters. | As a particular example, we compare the crudely predicted wobble to that observed in SS433 and find agreement to order of magnitude for a plausible choice of disc parameters. |
The stochastic nature of the tlt distinguishes it from that of an ordered precession We conclude in Sec. | The stochastic nature of the tilt distinguishes it from that of an ordered precession We conclude in Sec. |
5. | 5. |
Consider (he random walk on a two-dimensional square latlice of points 1943). | Consider the random walk on a two-dimensional square lattice of points \citep{ch43}. |
. We label a point on the lattice by coordinates (8.ᾧ). and (he spacing between adjacent lattice points bv A@, | We label a point on the lattice by coordinates $(\theta,\phi)$, and the spacing between adjacent lattice points by $\Delta \theta$. |
We also break time into discrete intervals A/. | We also break time into discrete intervals $\Delta t$. |
We use these coordinates in anticipation of our application to the random walk of a unit. vector tip on the surface of a sphere. locally approximated as a plane for small angles. | We use these coordinates in anticipation of our application to the random walk of a unit vector tip on the surface of a sphere, locally approximated as a plane for small angles. |
At time /— M. a point particle located at (8.0) moves bv time / to one ofthe 4 points | At time $t-\Delta t$ , a point particle located at $(\theta,\phi)$ moves by time $t$ to one ofthe 4 points |
The amount aud nature of Dark Matter present in the Universe is an nuportant question for cosmologv (sec ce. White et al. | The amount and nature of Dark Matter present in the Universe is an important question for cosmology (see e.g. White et al. |
for current status). | \nocite{White:1996} for current status). |
On galactic scales (Aslan 1992).. dynamical studies (Zaritsky as well as macrolensing analyses (Carollo et al. | On galactic scales (Ashman \nocite{Ashman:1992}, dynamical studies (Zaritsky \nocite{Zaritsky:1992} as well as macrolensing analyses (Carollo et al. |
show that np to of the ealactic masses might not ὃς visible. | \nocite{Carollo:1995} show that up to of the galactic masses might not be visible. |
One plausible explanation is that the stellar éoutent of ealaxies is embedded iu a dark halo. | One plausible explanation is that the stellar content of galaxies is embedded in a dark halo. |
Primordial ameleosynthesis: (Walker ct al. | Primordial nucleosynthesis (Walker et al. |
1991: Copi et al. | 1991; Copi et al. |
1995) predicts a larger nuuber of barvous thin what is secu (Persie Salucci 1992).. aud so dark baryous hidden iu | 1995) \nocite{Walker:1991,Copi:1995} predicts a larger number of baryons than what is seen (Persic Salucci \nocite{Persic:1992}, , and so dark baryons hidden in |
Given the low resolution of the sspectra we will restrict the analysis to fitting simple phenomenological spectral components to obtain the broad shape of the variability components. | Given the low resolution of the spectra we will restrict the analysis to fitting simple phenomenological spectral components to obtain the broad shape of the variability components. |
We start. by fitting the three Fourier. frequency spectra. of the combined orbits 9991003 with a power law. allected only by Galactic absorption (ng=L7.107" en). inthe 0.810 keV baud. eiving values for the power law slope of 1.934: 0.01 for LE. 1.85 c 0.03 for ME and 1.66 40.04 for LIE. | We start by fitting the three Fourier frequency spectra of the combined orbits 999–1003 with a power law, affected only by Galactic absorption $n_{\rm H} =1.7 \times 10 ^{20}$ $^2$ ), in the 0.8–10 keV band, giving values for the power law slope of $\pm$ 0.01 for LF, 1.85 $\pm$ 0.03 for MF and 1.66 $\pm$ 0.04 for HF. |
Phe extrapolation of these fits to lower energies shows a "soft excess’ over the power law in the three Fourier-frequeney spectra. | The extrapolation of these fits to lower energies shows a `soft excess' over the power law in the three Fourier-frequency spectra. |
A similar ellect is seen in the sspectrum of the combined orbit. 1000.1003. where the high energy power law slopes are 1.96 + 0.01 for LE. 1.94 x 0.02 or ME and 1.75 + 0.04 for H1E.(? cof = 3.8 for 33 degrees of freedom). as shown in Fig. 4.. | A similar effect is seen in the spectrum of the combined orbit 1000–1003, where the high energy power law slopes are 1.96 $\pm$ 0.01 for LF, 1.94 $\pm$ 0.02 for MF and 1.75 $\pm$ 0.04 for $\chi^2/$ dof = 3.8 for 33 degrees of freedom), as shown in Fig. \ref{rms_999_1003}. |
A similar shape was found in the FR spectrum analysis of LCG6-80-15. presented by Papaclakisetal.(2005). | A similar shape was found in the FR spectrum analysis of MCG–6-30-15, presented by \citet{Papadakis_frspec}. |
. The soft excess in the sspectra can be modelled as a broken power law or as a single power Law with an absorption feature around 0.7 keV. We itted the three frequency ranges with a broken power Law alfected by galactic absorption. | The soft excess in the spectra can be modelled as a broken power law or as a single power law with an absorption feature around 0.7 keV. We fitted the three frequency ranges with a broken power law affected by galactic absorption. |
Ehe resulting fit parameters are listecl in Table 1.. | The resulting fit parameters are listed in Table \ref{broken_po_table}. |
The low and high energv slopes get [latter with increasing frequency. | The low and high energy slopes get flatter with increasing frequency. |
The low energy slope changes by approximately 0.150.3 and the high energy slope changes by ~0.2. unlike the case of ALCG6-30-15. where the soft. power law showed no significant. frequency dependence (Papaclakisetal..2005). | The low energy slope changes by approximately 0.15–0.3 and the high energy slope changes by $\sim
0.2$, unlike the case of MCG–6-30-15, where the soft power law showed no significant frequency dependence \citep{Papadakis_frspec}. |
. The broken power law model gencrally does not produce a good fit to the sspectra of aand the residuals show svstematic structure as shown in the top panel of Fig. 5.. | The broken power law model generally does not produce a good fit to the spectra of and the residuals show systematic structure as shown in the top panel of Fig. \ref{edge_po_fit}. |
A better fit to the sspectra is obtained with a single power law allected by colc Galactic absorption and a warn absorber. | A better fit to the spectra is obtained with a single power law affected by cold Galactic absorption and a warm absorber. |
Given the [ow resolution of the spectra presented here. we initially simpA modeled the warm. absorber as two edges. lor O VILE anc O VILL using the Nspec mocel zedge. | Given the low resolution of the spectra presented here, we initially simply modeled the warm absorber as two edges, for O VII and O VIII, using the Xspec model zedge. |
The inclusion of the second absorption edge. however. does not improve the fits significantly. so we dropped one of the edge components anc used a free energy for the remaining edge instead. to simula | The inclusion of the second absorption edge, however, does not improve the fits significantly, so we dropped one of the edge components and used a free energy for the remaining edge instead, to simulate |
galaxies the end of the spiral or at least the end of its svmmetric part lies near corotation or inner 4/I-resonancoe. | galaxies the end of the spiral or at least the end of its symmetric part lies near corotation or inner 4/1-resonance. |
The possible independence of corotation resonance radius rom the bar size in the galaxies with smallest. bars could suggest that the spiral in these galaxies indeed has a cillerent xutern speed. whose corotation radius is regulated by the galaxy size. | The possible independence of corotation resonance radius from the bar size in the galaxies with smallest bars could suggest that the spiral in these galaxies indeed has a different pattern speed, whose corotation radius is regulated by the galaxy size. |
In Fig. | In Fig. |
9 we compare Roto parameters Qe and Qs from ?.. which characterize the strengths of bar and spiral components (maximum tangential force. divided: by he azimuthally averaged racial force in each radius). | \ref{coro_barstrength} we compare $\mathcal{R}$ to parameters $Q_B$ and $Q_S$ from \citet{buta2005}, which characterize the strengths of bar and spiral components (maximum tangential force divided by the azimuthally averaged radial force in each radius). |
With one exception. galaxies with slow bars CR. 1.4) have Qex0.3. | With one exception, galaxies with slow bars $\mathcal{R} > 1.4$ ) have $Q_B
\le 0.3$. |
However. there are several galaxies with similar xw strength. but fast bar. | However, there are several galaxies with similar bar strength but fast bar. |
Phe combined results with bar size and strength are in accordance with ? who found by morphological arguments that the flat. (and typically long) xws tend to be faster rotators than the exponential (ancl vpically short) bars. | The combined results with bar size and strength are in accordance with \citet{elmegreen95a} who found by morphological arguments that the flat (and typically long) bars tend to be faster rotators than the exponential (and typically short) bars. |
There is no clear correlation between xutern speed ancl spiral strength. | There is no clear correlation between pattern speed and spiral strength. |
The top frame of Fig. | The top frame of Fig. |
10. shows R as a function of bulee-to- Dux fraction (2/7 -ratio) in df-bane (seealso?).. | \ref{coro_bt_etc} shows $\mathcal{R}$ as a function of bulge-to-total flux fraction $B/T$ -ratio) in $H$ -band \citep[see also][]{salo2007}. |
The apparent bulge contribution is one of the eriteria in Llubble classification. so it is not surprising that we ect almost a mirror image of Fig. 6.. | The apparent bulge contribution is one of the criteria in Hubble classification, so it is not surprising that we get almost a mirror image of Fig. \ref{coro_trc}. |
We studied the dependeney between R anc absolute D- (calculated from the D-magnitudes in LCS) and distances by 2.. but found no clear correlation. | We studied the dependency between $\mathcal{R}$ and absolute $B$ (calculated from the $B$ -magnitudes in RC3) and distances by \citet{tully88}, but found no clear correlation. |
Phe same was the situation with absolute A'-magnitudes based on 2ALASS data (7)/ and neither there was correlation with colour. | The same was the situation with absolute $K$ -magnitudes based on 2MASS data \citep{jarrett2003} and neither there was correlation with colour. |
Llowever. if we plot the galaxies in (2/7.I )-plane (rig. 10.. | However, if we plot the galaxies in $(B/T,K)$ -plane (Fig. \ref{coro_bt_etc}, |
bottom). then the fast and slow galaxies roughly occupy different areas: slower bar favouring lower bulge fraction and/or higher galaxy luminosity. | bottom), then the fast and slow galaxies roughly occupy different areas: slower bar favouring lower bulge fraction and/or higher galaxy luminosity. |
According to gas dynamical simulations by 2.. the existence of leading offset. dust lanes in the bars tells us two things: that there is an ILR and that R=1240.2. | According to gas dynamical simulations by \citet{athanassoula92b}, the existence of leading offset dust lanes in the bars tells us two things: that there is an ILR and that $\mathcal{R}=1.2 \pm 0.2$. |
The first condition comes from the existence of so-called. w2-orbits. which require the presence of an LLR: the dust lanes form as the orientation of the orbits of gas clouds change from parallel orientation of wy-orbits (with respect to the bar) in the outer parts of the bar to the perpendieular orientation of e-orbits in the inner parts. | The first condition comes from the existence of so-called $x_2$ -orbits, which require the presence of an ILR: the dust lanes form as the orientation of the orbits of gas clouds change from parallel orientation of $x_1$ -orbits (with respect to the bar) in the outer parts of the bar to the perpendicular orientation of $x_2$ -orbits in the inner parts. |
The second condition comes from the gas cvnamical simulation. only in this parameter region the shape of the shock corresponds to the observed dust. morphology. | The second condition comes from the gas dynamical simulation, only in this parameter region the shape of the shock corresponds to the observed dust morphology. |
We have checked the dust. lane morphology of the sample galaxies from. OSUBSGS images ancl images [rom various other sources. NNED and the web page of de Vaucouleurs Atlas ofGalaxies (?).. | We have checked the dust lane morphology of the sample galaxies from OSUBSGS images and images from various other sources, NED and the web page of de Vaucouleurs Atlas of Galaxies \citep{buta2007a}. |
Omitting cases where the inclination makes recognition and classification of dust lanes ambiguous. we found that more than third of the galaxies in the sample does not show classic dust lane morphology. | Omitting cases where the inclination makes recognition and classification of dust lanes ambiguous, we found that more than third of the galaxies in the sample does not show classic dust lane morphology. |
ο. is a massive eclipsiug N-vav binary svsten at an advanced evolutionary stage. | SS433 is a massive eclipsing X-ray binary system at an advanced evolutionary stage. |
It is recognized as a superciitically accreting mücroquasar with a precessine accretion disk axd wildly relativistic (c~-- 0.266) jets. | It is recognized as a supercritically accreting microquasar with a precessing accretion disk and mildly relativistic $v\approx 0.26\,c$ ) jets. |
Since its discovery in 1975 (Clark awl πάσα 1975. Mareou et al. | Since its discovery in 1978 (Clark and Murdin 1978, Margon et al. |
1979). this unique ταν binary has been deeply investigaed in optical radio. ando N-ravs (see reviews by Maureon 1981. Cherepasheluκ 1988. 2002 and Fabrika 2001 for more detail and refererwes). | 1979), this unique X-ray binary has been deeply investigated in optical, radio, and X-rays (see reviews by Margon 1984, Cherepashchuk 1988, 2002 and Fabrika 2004 for more detail and references). |
Amoug a ozen known σοςi!asars. SS133 is distinguished by its unique properties. | Among a dozen known microquasars, SS433 is distinguished by its unique properties. |
1) Tn this svstem the optical star fIs its Roche lohe aud mass transfer occurs on theriual tine scale with the hugerate of AL~10.1A. per year aud accretion onto the relativistic object is supercritical (Shakura aud Suuvaev 1973). | 1) In this system the optical star fills its Roche lobe and mass transfer occurs on thermal time scale with the huge rate of $\dot M \sim
10^{-4} M_\odot$ per year and accretion onto the relativistic object is supercritical (Shakura and Sunyaev 1973). |
2) Two strouglv collimated (~ 17) oppositely directed jets enanate with a velocity of SO.000+1.00) lan per second from the ceuter of the accretion disk. | 2) Two strongly collimated $\sim 1^\circ$ ) oppositely directed jets emanate with a velocity of $80,000\pm 1,000$ km per second from the center of the accretion disk. |
The jets are observed not sporadically, as iu most nicroquasars. but yersist over tens of vears with virtually constant velocity. | The jets are observed not sporadically, as in most microquasars, but persist over tens of years with virtually constant velocity. |
3) The accretion disk aud relativistic jets regularly recess With a period of 162.5 davs. | 3) The accretion disk and relativistic jets regularly precess with a period of 162.5 days. |
Both the precession xxiod aud the disk inclination angle to the orbital plane (~20° ) on average remain constaut over tenus of vears. | Both the precession period and the disk inclination angle to the orbital plane $\sim 20^\circ$ ) on average remain constant over tens of years. |
1) The source exhibits orbital eclipsineo periodicity with 134.082. | 4) The source exhibits orbital eclipsing periodicity with period $P_{orb}=13^d. 082$. |
The shape of the optical Πο curveo varies significantly with- precessional phase o (Gorauskiy et al. | The shape of the optical light curve varies significantly with precessional phase (Goranskij et al. |
1998ab. Cherepaslhichiuk and Yarikov 1991). | 1998ab, Cherepashchuk and Yarikov 1991). |
The orbital period is found to be very stable over more than 25 vears. despite the high mass exchange rate between t16 binary components aud stroug wilxd nass loss (οzz2000 au/s Mz10LAL, 'vi) frou the supercritical accretion cisk. | The orbital period is found to be very stable over more than 25 years, despite the high mass exchange rate between the binary components and strong wind mass loss $v\approx 2000$ km/s, $\dot M\simeq 10^{-4} M_\odot$ /yr) from the supercritical accretion disk. |
The τιsolved puzzles of 88133 sil remaining to solved inchide: (1) the nature of the relativistic object. (2) the mechaisi of collimation and acceleration of matter in jets to fhe relativistic velocity ~&y000 ins. (3) t nature of 16 precessional phenomena in this X-ray binary system (CIjrabarti. 2002). | The unsolved puzzles of SS433 still remaining to be solved include: (1) the nature of the relativistic object, (2) the mechanism of collimation and acceleration of matter in jets to the relativistic velocity $\sim 80,000$ km/s, (3) the nature of the precessional phenomena in this X-ray binary system (Chakrabarti, 2002). |
The observations of SS133 ii 2003 discovered a hard (p to 100 keV) N-rav spectrum iu this mooreriticallv. accreting wicroquasar (Chereywashchuls et al. | The observations of SS433 in 2003 discovered a hard (up to 100 keV) X-ray spectrum in this supercritically accreting microquasar (Cherepashchuk et al. |
2003. 2001). suggesting the presence of an exteuded hot (with a temperature up to 1 Po) rosjou iu the central parts of the accreion disk. | 2003, 2004), suggesting the presence of an extended hot (with a temperature up to $10^8$ K) region in the central parts of the accretion disk. |
These new data made it possible to compare the eclipse characteristics of SS133 at different energies: soft X-rays (2-10 keV. the data). soft aud iecimm N-rav {1-27 keV. the data). hard Noavs (20-70 keV. the data). and the optical. | These new data made it possible to compare the eclipse characteristics of SS433 at different energies: soft X-rays (2-10 keV, the data), soft and medium X-ray (1-27 keV, the data), hard X-rays (20-70 keV, the data), and the optical. |
This comparison alows us fo investigate the iunermost structure of the supercritical accretion disk and to constrain the basic parameters of the bilary svstena | This comparison allows us to investigate the innermost structure of the supercritical accretion disk and to constrain the basic parameters of the binary system. |
This paper is organized as follows: Secjon 2 lists the participants of the observational campaign. | This paper is organized as follows: Section 2 lists the participants of the observational campaign. |
Section 3 describes the N-vav spectra and helt curveS obtained. | Section 3 describes the X-ray spectra and light curves obtained. |
Section L preseuts the analysis of lard X-ray oecessional and orbital variability of SS133. | Section 4 presents the analysis of hard X-ray precessional and orbital variability of SS433. |
Section 5 reports the results of optical and infrared photometry. | Section 5 reports the results of optical and infrared photometry. |
Section 6 describes in detail the simmitaneous optical sο | Section 6 describes in detail the simultaneous optical spectroscopy |
This appendix provides some additional notes to the methods emploved iu this paper. | This appendix provides some additional notes to the methods employed in this paper. |
For a more in-depth discussion of the topics presented here. please refer to the cited publications. | For a more in-depth discussion of the topics presented here, please refer to the cited publications. |
Iu Gaussian statistics. our probability densities are fully defined by the first and second statistical momeuts. be. their meas aud covariauces. | In Gaussian statistics, our probability densities are fully defined by the first and second statistical moments, i.e. their means and covariances. |
Two random vectors. 8jaud 84,4. are said to be uncorrelated when their covariauce (Csa, yas zero: where E|sj] is the expectation value of s; which can be approximated by the mean in this case by with AZ απο the number of data poiuts in the time series. | Two random vectors, $\bf{s_{l}}$and $\bf{s_{l+1}}$, are said to be uncorrelated when their covariance ${\bf C}_{\bf{s}_l, \bf{s}_{l+1}}$ ) is zero: where $E[{s_{l}}]$ is the expectation value of ${s_{l}}$ which can be approximated by the mean in this case by with $M$ being the number of data points in the time series. |
Furtheriiore. we define two raudonm variables 6s; auc s;) 1) to be orthogonal. when both their expectation values. in addition to their covariance are zero: We can abwavs fiud an affine. linear trausformation from a correlated set of variables to an orthogonal one. | Furthermore, we define two random variables ${s_{l}}$ and ${s_{l+1}}$ ) to be orthogonal, when both their expectation values, in addition to their covariance are zero: We can always find an affine, linear transformation from a correlated set of variables to an orthogonal one. |
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